Version 4.0 June 1996 Wide Field and Planetary Camera 2 Instrument Handbook SPACE TELESCOPE SCIENCE INSTITUTE Operated by the Association of Universities for Research in Astronomy, Inc., for the National Aeronautics and Space Administration User Support For prompt answers to questions, please contact the Science Support Division Help Desk. * E-mail: help@stsci.edu * Phone: (410) 338-1082 World Wide Web Information, software tools, and other resources are available on the WFPC2 World Wide Web page: * URL:http://www.stsci.edu/ftp/instrument_news/WFPC2/top_wfpc2.html WFPC2 Instrument Team Name Title Phone e-mail ------------------------------------------------------------------------------ Brad Whitmore Group Lead 410-338-4743 whitmore@stsci.edu John Biretta Technical Scientist 410-338-4917 biretta@stsci.edu Stefano Casertano Calibration Scientist 410-338-4752 stefano@stsci.edu Keith Noll User Support Scientist 410-338-1828 noll@stsci.edu Sylvia Baggett Cal. Pipeline Scientist 410-338-5054 sbaggett@stsci.edu Harry Ferguson Instrument Scientist 410-338-5098 ferguson@stsci.edu Andrew Fruchter Instrument Scientist 410-338-5018 fruchter@stsci.edu Christopher O'Dea Instrument Scientist 410-338-2590 odea@stsci.edu Massimo Stiavelli Instrument Scientist 410-338-4835 mstiavel@stsci.edu Anatoly Suchkov Instrument Scientist 410-338-4979 suchkov@stsci.edu Jean Surdej Instrument Scientist 410-338-4984 surdej@stsci.edu Shireen Gonzaga Data Analyst 410-338-4412 gonzaga@stsci.edu Inge Heyer Data Analyst 410-338-5017 heyer@stsci.edu Matt McMaster Data Analyst 410-338-4463 mcmaster@stsci.edu Michael S. Wiggs Data Analyst 410-338-4998 wiggs@stsci.edu J.-C. Hsu Scientific Programmer 410-338-4760 hsu@stsci.edu ------------------------------------------------------------------------------ The WFPC2 Investigation Definition Team: John T. Trauger, Christopher J. Burrows, John Clarke, David Crisp, John Gallagher, Richard E. Griffiths, J. Jeff Hester, John Hoessel, Jon Holtzman, Jeremy Mould, and James A. Westphal Handbook Authors and Contributors: John Biretta, Chris Burrows, Jon Holtzman, Inge Heyer, Mark Stevens, Sylvia Baggett, Stefano Casertano, Mark Clampin, Andrew Fruchter, Harry Ferguson, Ron Gilliland, Richard Griffiths, John Krist, Keith Noll, Christopher O'Dea, Massimo Stiavelli, Anatoly Suchkov, Jean Surdej, and Brad Whitmore Revision History Instrument Version Date Editor ------------------------------------------------------------------------------ WF/PC-1 1.0;2.0;2.1 October 1985; May 1989; May 1990 Richard Griffiths WF/PC-1 3.0 April 1992 John W. MacKenty WFPC2 1.0;2.0;3.0 March 1993; May 1994; June 1995 Christopher Burrows WFPC2 4.0 June 1996 John A. Biretta ------------------------------------------------------------------------------ In publications please refer to this document as "Biretta, J. A., et al. 1996, WFPC2 Instrument Handbook, Version 4.0 (Baltimore: STScI)." @ 1996 by STScI. Table of Contents CHAPTER 1: Introduction Instrument Overview Field-of-View Spectral Filters Quantum Efficiency and Exposure Limits CCD Detector Technology UV Imaging Aberration Correction and Optical Alignment Which Instrument to Use: WFPC2, FOC, NICMOS, or STIS? Comparison of WFPC2 and FOC Comparison of WFPC2 and NICMOS Comparison of WFPC2 and STIS History of WFPC2 The Previous vs. Current Generation: WFIPC-1 vs. WFPC2 Organization of this Handbook What's New in Version 4.0 WFPC2 Handbook on the WWW The Help Desk at STScI Further Information CHAPTER 2: Instrument Description Science Objectives WFPC2 Configuration, Field-of-View, and Resolution Overall Instrument Description Quantum Efficiency Shutter Serial Clocks Overhead Times CCD Orientation and Readout Calibration Channel CHAPTER 3: Optical Filters Introduction Choice of Broad Band Filters Linear Ramp Filters Spectral Response Target Locations LRF Photometric Calibration Redshifted [OII] Quad Filters Polarizer Quad Filter Methane Quad Filter Wood's Filters Red Leaks in UV Filters Apertures CHAPTER 4: CCD Performance Introduction Quantum Efficiency Dynamic Range Bright Object Artifacts Blooming Horizontal Smearing Diff raction Effects and Ghost Images Residual Image Quantum Efficiency Hysteresis Flat Field Response Dark Backgrounds Sources of Dark Current Darktime Cosmic Rays Radiation Damage and Hot Pixels Charge Transfer Efficiency Read Noise and Gain Settings CHAPTER 5: Point Spread Function Effects of OTA Spherical Aberration Aberration Correction Wavefront Quality CCD Pixel Response Function Model PSFs PSF Variations with Field Position PSF Variations with Time Large Angle Scattering Ghost Images Optical Distortion CHAPTER 6: System Throughput and SNR I Exposure Time Estimation System Throughput On-Line Exposure Time Calculator Target Count Rates Count Rates for Stellar Sources Count Rates for Power Law Sources Count Rates for Emission Line Sources Sky Background Signal-to-Noise Ratio Estimation Point Sources--PSF Fitting Point Sources--Aperture Photometry Extended Sources Exposure Time Estimation Sample SNR Calculations 1 Point Sources Extended Sources Emission Line Sources Red Leaks in UV Filters Time Dependence of UV Response CHAPTER 7: Observation Strategies Observing Faint Targets Observing Bright Targets Observing Faint Targets Near Bright Objects Cosmic Rays 151 Choosing Exposure Times Dither Strategies 155 Pointing Accuracy Absolute Pointing Accuracy Updates to Aperture / Coordinate Systems Pointing Repeatability Tracking Modes 159 CCD Position and Orientation on Sky Polarization Observations Observing with Linear Ramp Filters Emission Line Observations of Galaxy Nuclei CHAPTER 8: Calibration and Data Reduction Calibration Observations and Reference Data Flat Fields Dark Frames Bias Frames Data Reduction and Data Products Pipeline Processing Fluxes and Standard Magnitudes Color Transformations - Primary Photometric Filters Cycle 5 Calibration Plan Cycle 6 Calibration Plan CHAPTER 9: References References Instrument Science Reports APPENDIX A: Passband Plots Passbands for each Filter in Isolation Passbands including the System Response Normalized Passbands including the System Re-sponse APPENDIX B: Point Source SNR Plots Plots for Estimating Point Source SNR 223 APPENDIX C: Acronyms INDEX List of Figures Figure 1.1: WFPC2 Field-of-View Projected on the Sky Figure 2.l: Wide Field Planetary Camera Concept Illustration Figure 2.2: WFPC2 Optical Configuration Figure 2.3: Cooled Sensor Assembly Figure 2.4: WFPC2 + OTA System Throughput Figure 3.1: Summary of Normalized Filter Curves Figure 3.2: Ramp Filter Peak Transmission Figure 3.3: Ramp Filter Dimensionless Widths Figure 3.4: FR418N and FR533N Wavelength Mapping Figure 3.5: FR68ON and FR868N Wavelength Mapping Figure 3.6: Polarizer Quads Figure 3.7: Polarizer Transmission Figure 3.8: Methane Quad Filter Figure 3.9: Wood's Filters Figure 3.10: UV Filter Red Leaks Figure 3.11: Precise CCD Alignments and Primary Aperture Locations Figure 4.l: MPP Operating Principle Figure 4.2: Pre-flight DQE Measurements on WFPC2 CCDs Figure 4.3: Saturated Stellar Image Showing Horizontal Smearing Figure 4.4: WFPC2 CCD Flat Field Figure 4.5: Average Dark Rates vs. CCD Row Figure 4.6: Dark Signal vs. Cosmic Ray Flux Figure 4.7: Histogram of Cosmic Ray Event Sizes Figure 4.8: Comparison of Star Images and Cosmic Ray Events Figure 4.9: Histogram of Cosmic Ray Event Energies Figure 4.10: Hot Pixel Histogram Figure 4.11: Images Illustrating CTE Residual Trail Figure 5.l: PSF Surface Brightness Figure 5.2: Encircled Energy Figure 5.3: Encircled Energy for CCD PCl Figure 5.4: Encircled Energy for CCD WF3 Figure 5.5: PSF Variations with Field Position - Aberrations Figure 5.6: PSF Variations with Field Position - Obscuration Shifts Figure 5.7: Measured OTA Focus Position (microns) vs. Time Figure 5.8: Measured Aperture Correction, V(r) - V(r=10 pix) Figure 5.9: Large Angle Scattering Figure 5.10: Saturated Stellar Image Showing Filter Ghosts Figure 5.11: Saturated Stellar Image Showing Field Flattener Ghost on WF2 Figure 5.12: Integrated Photometry Correction Induced by Camera Distortions Figure 6.1: Giant Elliptical Galaxy Figure 6.2: Sample Fill-out Form for WFPC2 On-Line Exposure Time Calculator Figure 6.3: Sample Results from WFPC2 ETC - Point Source Figure 6.4: Sample Results on CR-SPLITing from WFPC2 ETC Figure 6.5: Point Source + Stellar Background Fill-out Form for WFPC2 ETC Figure 6.6: Sample Output from WFPC2 ETC - Point Source + Background Figure 6.7: Extended Source Form for WFPC2 ETC Figure 6.8: Sample Output from WFPC2 ETC - Extended Source Figure 6.9: Point Source Form for WFPC2 ETC - Emission Line & LRF Figure 6.10: Sample Output for WFPC2 ETC - Emission Line & LRF Figure 6.11: Photometric Monitoring Data for WFPC2 Figure 6.12: Throughput for the Fl 7OW Filter Following Decontaminations Figure 6.13: Change in Throughput vs. Time Figure 7.1: Example of Scattered Earth Light Figure 7.2: Impact of OTA Focus Shift on PSF Subtraction Figure 7.3: Bright Object Avoidance Regions Near WFPC2 FOV Figure 7.4: Example of PC1 "Direct" Stray Light Ghost Figure 7.5: Example of PC1 "Diffraction" Stray Light Ghost Figure 7.6: Event Timings During a 60s WFPC2 Exposure Figure 7.7: Event Timings During a 1 00s WFPC2 Exposure Figure 7.8: ORIENT Definition, Aperture Positions, and CCD Alignments Figure 7.9: Example of ORIENT and POS TARG Selection Figure 7.10: Example of ORIENT and POS TARG Selection Figure 8.1: F336W-F439W against Johnson U-B Figure 8.2: F439W-F555W against Johnson B-V Figure 8.3: F555W-F814W against Johnson V-lc Figure 8.4: F555W-F675W against Johnson V-Rc Figure 8.5: F675W-F814W against Cousins RC-IC Figure B.1: Point Source SNR vs. V+AB, for Fl 6OBW Filter Figure B.2: Point Source SNR vs. V+ABV for F218W Filter Figure B.3: Point Source SNR vs. V+ABV for F255W Filter Figure B.4: Point Source SNR vs. V+ABV for F30OW Filter Figure B.5: Point Source SNR vs. V+ABV for F336W Filter Figure B.6: Point Source SNR vs. V+ABV for F41 OM Filter Figure B.7: Point Source SNR vs. V+ABV for F439W Filter Figure B.8: Point Source SNR vs. V+ABV for F502N Filter Figure B.9: Point Source SNR vs. V+ABV for F547M Filter FigureB.10: PointSourceSNRvs.V+ABvforF555WFilter FigureB.11: PointSourceSNRvs.V+ABvforF606WFilter FigureB.12: PointSourceSNRvs.V+ABvforF675WFilter FigureB.13: PointSourceSNRvs.V+ABvforF702WFilter FigureB.14: PointSourceSNRvs.V+ABvforF814WFilter List of Tables Table 1.1: Comparison of WFPC2, FOC, NICMOS, and STISI Table 1.2: Comparison of WFPC2 and FOC Detective Efficiencies Table 1.3: Comparison of WFPC2 and NICMOS Count Rates Table 1.4: Comparison of WFPC2 and STIS Detective Efficiencies Table 2.1: Summary of Camera Format Table 2.2: WFPC2 Dynamic Range in a Single Exposure Table 2.3: Quantized Exposure Times (Seconds) Table 2.4: Instrument Overheads Table 2.5: Inner Field Edges Table 3.1: WFPC2 Simple Filter Set Table 3.2: WFPC2 Quad and Ramp Filters Table 3.3: Ramp Filter FR418N Parameters Table 3.4: Ramp Filter FR533N Parameters Table 3.5: Ramp Filter FR68ON Parameters Table 3.6: Ramp Filter FR868N Parameters Table 3.7: Aperture Locations and Wavelengths for Ramp Filters Table 3.8: Vignetted Wavelengths for Ramp Filters Table 3.9: Redshifted [011] Quad Filter Elements Table 3.10: Polarizer Quad Filter Table 3.11: Methane Band Quad Filter Table 3.12: Wood's Filters Table 3.13: Red Leak in UV Filters Table 3.14: Aperture Definitions Table 3.15: Updates to (V2,V3) Positions of WFPC2 Apertures Table 4.1: Comparison of WF/PC-1 and WFPC2 CCDs Table 4.2: Dark Count Rates Table 4.3: Signal Chain Gains Table 5.1: Wavef ront Error Budget Table 5.2: Aberrations in Each Camera Table 5.3: PC Point Spread Functions Table 5.4: WFC Point Spread Functions Table 5.5: Cubic Distortion Coefficients Table 6.1: System Efficiencies and Zeropoints Table 6.2: ABV as a Function of Wavelength Table 6.3: Sky Brightness Table 6.4: Sky Count Rate per Pixel (Psky) Table 6.5: Sharpness Table 6.6: Parameters for Pt. Source SNR Est. - PSF Fitting Table 6.7: Encircled Energy for Representative Filters Table 6.8: Parameters for Pt. Source SNR Est. - Aperture Photom Table 6.9: Parameters for Extended Source SNR Estimation Table 6.10: Change in WFPC2 Throughput Over 30 Days Table 7.1: O'RIENTs for Avoiding Bloom and Diffraction Spikes Table 7.2: Approx. PSF Subtraction Artifact Magnitudes Table 7.3: Recommended Exposure Splittings Table 7.4: Basic Time to Execute Single Non-CR-SPLIT Exposure Table 8.1: Conversion from STMAG to UBVRI and Cousins RI Table 8.2: Summary of Cycle 5 Calibration Plans Table 8.3: Summary of Cycle 6 Calibration Plans Table B.l: ABV as a Function of Wavelength Table C.1: Acronyms Introduction In This Chapter... Instrument Overview Which Instrument to Use: WFPC2, FOC, NICMOS, or STIS? History of WFPC2 The Previous vs. Current Generation: WF/PC-1 vs. WFPC2 Organization of this Handbook 15 What's New in Version 4.0 WFPC2 Handbook on the WWW The Help Desk at STScI Further Information Instrument Overview Wide Field and Planetary Camera 2 (WFPC2) is a two-dimensional imaging photometer which is located at the center of the Hubble Space Telescope (HST) focal plane and covers the spectral range between approximately 1150A to 10500A. It simultaneously images a 150" x 150" "L"-shaped region with a spatial sampling of 0.1" per pixel, and a smaller 34" x 34" square field with 0.046" per pixel. The total system quantum efficiency (WFPC2+HST) ranges from 5 percent to 13 percent at visual wavelengths, and drops to ~0.5 percent in the far UV. Detection of faint targets will be limited by either the sky background (for broad filters) or by noise in the read-out electronics (for narrow and UV filters) with an RMS equivalent to 5 detected photons. Bright targets can cause saturation (>53000 detected photons per pixel), but there are no related safety issues. The sections below give a more detailed overview. 1.1.1 Field-of-View The WFPC2 field-of-view is divided into four cameras by a four-faceted pyramid mirror near the HST focal plane. Each of the four cameras contains an 8OOx8OO pixel Loral CCD detector. Three cameras operate at an image scale of 0.1" per pixel (F/12.9) and comprise the Wide Field Camera (WFC) with an "L" shaped field-of-view. The fourth camera operates at 0.046" per pixel (F/28.3) and is referred to as the Planetary Camera (PC). There are thus four sets of relay optics and CCD sensors in WFPC2. The four cameras are called PC1, WF2, WF3, and WF4, and their fields-of-view are illustrated in Figure 1. I (see also "CCD Position and Orientation on Sky" on page 159). Each image is a mosaic of three F/12.9 images and one F/28.3 image. Figure 1.1: WFPC2 Field-of-View Projected on the Sky. The readout direction is marked with arrows near the start of the first row in each CCD. The X-Y coordinate directions are for POS-TARG commands. The position angle of V3 varies with pointing direction and observation epoch, and is given in the calibrated science header by keyword PA_V3. 1.1.2 Spectral Filters The WFPC2 contains 48 filters mounted in 12 wheels of the Selectable Optical Filter Assembly (SOFA). These include a set of broad band filters approximating Johnson-Cousins UBVRI, as well as a set of wide U, B, V, and R filters, and a set of medium bandwidth Stromgren u, v, b, and y filters. Narrow band filters include those for emission lines of Ne V (3426A), CN (~3900A), [0III] (4363A and 5007A), He II (4686A), Hbeta (4861A), He I (5876A), [OI] (6300A), Halpha (6563A), [NII] (6583A), [SII] (6716A and 6731A), and [SIII] (9531A). The narrow-band filters are designed to have the same dimensionless bandpass profile. Center wavelengths and profiles are uniformly accurate over the filter clear apertures, and laboratory calibrations include profiles, blocking, and temperature shift coefficients. There are also two narrow band "quad" filters, each containing four separate filters which image a limited field-of-view: the UV quad contains filters for observing redshifted [OII] emission and are centered at 3767A, 3831A, 3915A, and 3993A. The Methane quad contains filters at 5433A, 6193A, 7274A, and 8929A. Finally, there is a set of narrow band "linear ramp filters" (LRFS) which are continuously tunable from 3170A to 9762A; these provide a limited field-of-view with diameter ~10". At ultraviolet wavelengths there is a solar-blind Wood's UV filter (1200-1900A). The UV capability is also enhanced by control of UV absorbing molecular contamination, the capability to remove UV absorbing accumulations on cold CCD windows without disrupting the CCD quantum efficiencies and flat field calibrations, and an internal source of UV reference flat field images. Finally, there is a set of four polarizers set at four different angles, which can be used in conjunction with other filters for polarimetric measurements. However, due to the relatively high instrumental polarization of WFPC2, they are probably suitable only for measurements on strongly polarized sources (> 3 percent polarized). 1.1.3 Quantum Efficiency and Exposure Limits The quantum efficiency (QE) of WFPC2+HST peaks at 13 percent in the red, and remains above 5 percent over the visible spectrum. The UV response extends to Lyman alpha wavelengths (QE~0.5 percent). Internal optics provide spherical aberration correction. Exposures of bright targets are limited by saturation effects, which appear above ~53000 detected photons per pixel (for setting ATD-GAIN= 15), and by the shortest exposure time which is 0.11 seconds. There are no instrument safety issues associated with bright targets. Detection of faint targets is limited by the sky background for broad band filters at visual wavelengths. For narrow band and ultraviolet filters, detections are limited by noise in the read-out amplifier ("read noise"), which contributes an RMS noise equivalent to ~5 detected photons per pixel. 1.1.4 CCD Detector Technology The WFPC2 CCDs are thick, front-side illuminated devices made by Loral. They support multi-pinned phase (MPP) operation which eliminates quantum efficiency hysteresis. They have a Lumogen phosphor coating to give UV sensitivity. Details may be summarized as follows: * Read noise: WFPC2 CCDs have ~5e- RMS read noise which provides good faint object and UV imaging capabilities. * Dark noise: Inverted phase operation yields low dark noise for WFPC2 CCDS. They are being operated at -88 degrees C and the median dark current is about 0.0045 e- pixel^-1 s^-1. * Flat field: WFPC2 CCDs have a uniform pixel-to-pixel response (< 2 percent pixel-to-pixel non-uniformity) which facilitates accurate photometric calibration. * CTE: Low level charge traps are present in the WFPC2 devices at the present operating temperature of -88'C. On low background star fields, the traps result in a loss of about 4 percent of the signal when a star image is clocked down through all rows of the CCD. In the presence of background, the effect is reduced. For most applications, CTE is negligible or calibratable and pre-flash exposures are not required. This avoids the increase in background noise, and the decrease in operational efficiency that results from a preflash. * Gain switch: Two CCD gains are available with WFPC2, a 7 e- DN^-1 channel which saturates at about 27000 e- (4096 DN with a bias of about 300 DN) and a 14 e- DN^-1 channel which saturates at about 53000 e-. The Loral devices have a full well capacity of -90,000 e- and are linear up to 4096 DN in both channels. * DQE: The peak CCD DQE in the optical is 40 percent at 7000A. In the UV (1100-4000A) the DQE is determined by the phosphorescent Lumogen coating, and is 10 - 15 percent. * Image Purge: The residual image resulting from a 100x (or more) full-well over-exposure is well below the read noise within 30 minutes. No CCD image purge is needed after observations of very bright objects. The Loral devices bleed almost exclusively along the columns. * Quantization: The systematic Analog-to-Digital converter errors have been largely eliminated, contributing to a lower effective read noise. * QEH: Quantum Efficiency Hysteresis (QEH) is not a significant problem in the Loral CCDs because they are frontside illuminated and use MPP operation. The absence of any significant QEH means that the devices do not need to be UV-flooded and the chips can be warmed monthly for decontamination purposes without needing to maintain a UV-flood. * Detector MTF: The Loral devices do suffer from low level detector MTF perhaps caused by scattering in the frontside electrode structure. The effect is to blur images and decrease the limiting magnitude by about 0.5 magnitudes. 1.1.5 UV Imaging WFPC2 had a design goal of 1 percent photometric stability at 1470A over a month. This requires a contamination collection rate of less than 47 ng cm^-2 month^-1 on the cold CCD window. Hence, the following features were designed into WFPC2 in an effort to reduce contaminants: 1. Venting and baffling, particularly of the electronics, were redesigned to isolate the optical cavity. 2. There was an extensive component selection and bake-out program, and specialized cleaning procedures. 3. Molecular absorbers (Zeolite) were incorporated. The CCDs were initially operated at -77 degrees C after launch, which was a compromise between being as warm as possible for contamination reasons, while being sufficiently cold for an adequate dark rate. However, at this temperature significant photometric errors were introduced by low-level traps in the CCDS. This problem with the charge transfer efficiency of the CCDs has been reduced since 23 April 1994 by operating the CCDs at -88 degrees C, but this leads to significantly higher contamination rates than hoped for. On-orbit measurements indicate that there is now a decrease in throughput at a repeatable rate of ~30 percent per month at 1700A, and moreover, that monthly decontamination procedures are able to remove the contaminants completely and recover this loss. 1.1.6 Aberration Correction and Optical Alignment WFTPC2 contains internal corrections for the spherical aberration of the HST primary mirror. These corrections are made by highly aspheric surfaces figured onto the Cassegrain relay secondary mirror inside each of the four cameras. Complete correction of the aberration depends on a precise alignment between the OTA pupil and these relay mirrors. Mechanisms inside WFPC2 allow optical alignment on-orbit. The 47 degree pick-off mirror has two-axis tilt capabilities provided by stepper motors and flexure linkages, to compensate for uncertainties in our knowledge of HST's latch positions (i.e., instrument tilt with respect to the HST optical axis). These latch uncertainties would be insignificant in an unaberrated telescope, but must be compensated for in a corrective optical system. In addition, three of the four fold mirrors, internal to the WFPC2 optical bench, have limited two-axis tilt motions provided by electrostrictive ceramic actuators and invar flexure mountings. Fold mirrors for the PC1, WF3, and WF4 cameras are articulated, while the WF2 fold mirror has a fixed invar mounting. A combination of the pick-off mirror and actuated fold mirror (AFMS) has allowed us to correct for pupil image misalignments in all four cameras. Since the initial alignment, stability has been such that mirror adjustments have not been necessary. The mechanisms are not available for GO commanding. 1.2 Which Instrument to Use: WFPC2, FOC, NICMOS, or STIS? In this section we compare briefly the performance of HST instruments with imaging capability in the UV to near-IR spectral range. As of this writing, both WFPC2 and FOC have capabilities in this area. Two new instruments, the Near-Infrared Camera and Multi-Object Spectrograph (NICMOS) and the Space Telescope Imaging Spectrograph (STIS) should be installed in HST during the second service mission in Early 1997. Important imaging parameters for all instruments are summarized in Table 1.1 below. Table 1.1: Comparison of WFPC2, FOC, NICMOS, and STIS Instrumental Imaging Parameters. 1.2.1 Comparison of WFPC2 and FOC Advantages of each instrument may be summarized as follows. WFPC2 advantages are: * Wider field-of-view, 150" x 150' vs. 7" x 7" or less. * Higher throughput at lambda > 4500A. * Better flat field accuracy: WFPC2 accuracy of ~1 percent vs. few percent for FOC. WFPC2 has no bright target safety issues, and can give useful data on faint targets near very bright objects. FOC can be damaged by bright objects. * Better short-term geometric stability: FOC is impacted by geometric changes during turn-on and as a function of orbit position. FOC advantages are: * Higher throughput at lambda < 45ooA. * Less impacted by far-UV contaminants. * PSF sampling: FOC offers 0.014" pixels vs. 0.0455" on WFPC2. The choice of instrument will be largely based on wavelength, required field-of-view, and need for good PSF sampling at short wavelengths. We note also that the FOC polarization data should be easier to calibrate, since there are fewer large-angle reflections than WFPC2, and since all polarization angles are available at all positions in the field-of-view. Table 1.2: Comparison of WFPC2 and FOC Detective Efficiencies. Instrument Filter Mean Bandpass Peak QE^b Wave- FWHM length(A) (A)^a ------------------------------------------------------------------------------ WFPC2 F122M 1420 100 0.12 percent FOC F14OW 1370 298 0.25 percent WFPC2 F255W 2586 393 0.45 percent FOC F22OW 2280 480 1.2 percent WFPC2 F336W 3342 500 2.9 percent FOC F342W 3410 702 5.4 percent WFPC2 F439W 4300 700 3.4 percent FOC F43OW 3940 832 4.7 percent WFPC2 F547M 5476 660 10.6 percent FOC F550M 5460 188 0.83 percent ------------------------------------------------------------------------------ a. Note that definition of FWHM is different from "effective width" elsewhere herein. b. Includes instrument and OTA. 1.2.2 Comparison of WFPC2 and NICMOS Both WFPC2 and NICMOS are capable of imaging at wavelengths between ~8000A and ~11,000A. At longer wavelengths NICMOS must be used; at shorter wavelengths either WFPC2, FOC, or STIS must be used. Table 1.3 below compares the detective efficiency of WFPC2 and NICMOS in the wavelength region where both instruments are viable. Count rates for a V=20 star of spectral class A0 are given for all filters at common wavelengths; the signal-to-noise (S/N) is also given for a 1 hour exposure of this same star. For bright continuum sources WFPC2 and NICMOS offer similar efficiency over the spectral range from 8800A to 10,50OA; the choice of instrument will likely depend on other factors such as field size and details of the passband shape. However, for very faint sources, the lower read noise of WFPC2 (5e- for WFPC2 vs. 30e- for NICMOS) should prove advantageous. Both instruments have a polarimetry capability, but the WFPC2 polarizers are not viable above 800OA; above this wavelength NICMOS must be used for polarimetry. Table 1.3: Comparison of WFPC2 and NICMOS Count Rates for V=20 AO Star. Instrument Filter Mean Effective Count Rate SNR in 1 Wavelength Width (e- s^-1) hour^a (A) (A) ------------------------------------------------------------------------------ WFPC2 F785LP 9,366 2095 14 215 F791W 8,006 1304 30 314 F814W 8,269 1758 33 333 F850LP 9,703 1670 7.1 150 FQCH4N (Quad D) 8,929 64 0.47 34,29^b F953N 9,546 52 0.21 19,15^b F1042M 10,443 611 0.20 18,15^b LRF^c 8,000 105 1.5 66 9,000 113 0.64 40 9,762 126 0.23 20 NICMOS F09OM^d 8,970 1885 2.0 90 F095N^d 9,536 88 0.14 12 F097N^d 9,715 94 0.18 16 F108N^d 10,816 94 0.17 15 F110W (Camera 1) 11,022 5920 6.6 160 F110W (Camera 2) 11,035 5915 14 260 F110W (Camera 3) 11,035 5915 26 350 ------------------------------------------------------------------------------ a. WFPC2 SNR assuming two 1800s exposures for cosmic ray removal. NICMOS SNR for central pixel of PSF. b. Values given for WFC (0.10" pixels) and PC (0.046" pixels). c. LRF filter is continuously tunable from 3710A to 9762A. LRF field-of-view is 10" x 10". d. These NICMOS filters are available only on Camera 1 which has 11" x 11" field-of-view. 1.2.3 Comparison of WFPC2 and STIS Both WFPC2 and STIS are capable of imaging over the same wavelength ranges between ~1150A and ~11000A. At much longer wavelengths NICMOS must be used. Advantages of each instrument may be summarized as follows. WFPC2 advantages are: * Wider field-of-view, 150" x 150" vs. 50" x 50" or less. * Greater selection of filters, including polarizers. * Flat field accuracy: WFPC2 is likely to have better flat fielding, since geometry is highly stable. STIS filters are near focal plane and geometry may be unstable due to Mode Selection Mechanism non-repeatability. * Field uniformity: STIS focus varies across field in image modes. * Bright Targets: WFPC2 has no bright target safety issues, and can give useful data on faint targets near very bright objects. STIS MAMAs can be damaged by bright objects. STIS advantages are: * Much higher UV throughput. * True solar blind imaging in UV due to MAMA detectors. WFPC2 CCDs are very sensitive to filter red-leak. * PSF sampling: STIS offers 0.024" pixels vs. 0.0455" on WFPC2. * High time resolution is possible ( tau ~125 micro seconds ) with the MAMAs detectors. Also the STIS CCD may be cycled on ~10s timescale using a sub-array. In general, WFPC2 has a much greater selection of filters and wider field-of-view than STIS, but STIS will have greater detective efficiency in the UV and for its long-pass and unfiltered modes. Table 1.2 below compares the detective efficiency for WFPC2 and STIS filters with similar bandpasses. For UV imaging STIS will be greatly superior due to higher throughput and insensitivity to filter red-leak; only if some detail of a WFPC2 filter bandpass were needed, would it be a viable choice. For both [OII] 3727A and [OIII] 5007A imaging STIS has much higher QE and will be preferred, unless the larger WFPC2 field-of-view is an important factor. The WFPC2 [OIII] filter is wider than its STIS counter-part, which may also be useful for redshifted lines. For broad-band imaging the unfiltered and 5500A long-pass modes of STIS again will have higher efficiency than WFPC2, though with reduced field-of-view. Table 1.4: Comparison of WFPC2 and STIS Detective Efficiencies. Instrument Filter Mean Bandpass Peak QE^b Wave- FWHM length (A) (A)^a ------------------------------------------------------------------------------ WFPC2 F122M 1420 100 0.12 percent STIS F25LYA 1216 85 4.0 percent WFPC2 F16OBW 1492 500 0.065 percent STIS FUV-MAMA ~1300 300 4.5 percent WFPC2 F255W 2586 393 0.45 percent STIS NUV-MAMA ~2400 1300 3.4 percent WFPC2 F375N 3738 42 0.8 percent STIS F28X50OII 3740 80 5 percent WFPC2 F502N 5013 37 5 percent STIS F28X50OIII 5007 5 15 percent WFPC2 F606W 5935 2200 13 percent STIS F28X50LP ~7300^c 2600 21 percent STIS F50CCD ~5800 4500 22 percent ------------------------------------------------------------------------------ a. Note that definition of FWHM is different from "effective width" elsewhere herein. b. Includes instrument and OTA. c. 5500A long pass filter. 1.3 History of WFPC2 The original Wide Field and Planetary Camera (WF/PC-1) served as the prototype for WFFPC2. In many respects the two instruments are very similar. Both were designed to operate from 1150A to 11000A, both use 8OO x 8OO CCD detectors, and both provide spatial samplings of ~0.045" and ~0.1 " per pixel. The development and construction of WF/PC-1 was led by Prof. J. A. Westphal, Principal Investigator (PI) of the California Institute of Technology. The instrument was built at the Jet Propulsion Laboratory (JPL) and was launched aboard HST in April 1990. It obtained scientific data until it was replaced by WFPC2 during the first service mission in December 1993. Because of its important role in the overall HST mission, NASA decided to build a second Wide Field and Planetary Camera (WFPC2) as a backup clone of WF/PC-1 even before HST was launched. WFPC2 was already in the early stages of construction at JPL when HST was launched. After the discovery of spherical aberration in the HST primary mirror, it was quickly realized that a modification of the WFPC2 internal optics could correct the aberration and restore most of the originally expected imaging performance. As a result, development of WFPC2 was accelerated. Dr. J. T. Trauger of JPL is the project PI for WFPC2 and led the Investigation Definition Team (IDT). The WFPC2 completed system level thermal vacuum (SLTV) testing at JPL in April and May 1993. Between June and November there were payload compatibility checks at Goddard Space Flight Center (GSFC), and payload integration at Kennedy Space Center (KSC). WF/PC-1 was replaced by WFPC2 during the first Servicing Mission in December 1993. WFPC2 was shown to meet most of its engineering and scientific performance requirements by testing conducted during the three month Servicing Mission Observatory Verification (SMOV) period following the servicing mission. The General Observer community has had access to WFPC2 since the start of Cycle 4 in January 1994. WFPC2 accurately corrects the HST spherical aberration, is a scientifically capable camera configured for reliable operation in space without maintenance, and is an instrument which can be calibrated and maintained without excessive operational overhead. It incorporates evolutionary improvements in photometric imaging capabilities. The CCD sensors, signal chain electronics, filter set, UV performance, internal calibrations, and operational efficiency have all been improved through new technologies and lessons learned from WF/PC-1 operations and the HST experience since launch. WFPC2 SMOV requirements were developed by the IDT, GSFC, and the STScI to include: verification of the baseline instrument performance; an optical adjustment by focusing and aligning to minimize coma; the estimation of residual wavefront errors from the analysis of star images; a photometric calibration with a core set of filters (including both visible and UV wavelengths); and the evaluation of photometric accuracy and stability over the full field with the core filter set. The results of these studies are documented in Holtzman, et al., 1995a and 1995b, and are summarized in this Handbook. Despite these successes, the first years of scientific operation of WFPC2 have revealed a number of relatively minor instrumental defects that were not expected from the pre-launch testing. These include a low-level charge transfer inefficiency, a higher than expected level of scattered light around bright objects, and variable and lower than expected ultraviolet (UV) efficiency. In addition, we have come to understand the instrument more fully -- particularly in terms of its overall photometric performance, geometric distortion, scale and alignments, hot pixels, and CCD traps. All of this new information is described here. 1.4 The Previous vs. Current Generation: WF/PC-1 vs. WFPC2 For historical reasons, it is useful to offer comparisons between the current WFPC2, and its predecessor WF/PC-1, which was returned to Earth in December 1993. * Field format: WF/PC-1 contained 8 cameras and CCDS, each CCD having 800 x 800 pixels. Four were used in the Planetary Camera mode (0.043" pixels), and four were used in the Wide Field Camera mode (0.10" pixels). The two pixel formats were selected by rotating the pyramid mirror by 45 degrees. WFPC2 budget and schedule constraints forced a reduction from 8 to 4 camera channels in August 1991. WFPC2 contains only 4 CCDS; the pyramid mirror is fixed and the 4 cameras are physically located in the bays occupied by the WF/PC-1 WFC. * Aberration correction: WF/PC-1 contained no correction for spherical aberration of the OTA primary mirror. Only about 15 percent of light from a stellar target fell into the core of the PSF (diameter -0.1"). WFPC2 incorporates corrective figures on the Cassegrain secondary mirrors inside the relay cameras, and as a result places ~60 percent of the light from a star inside a diameter of 0.1". Precise alignment of the OTA pupil on these mirrors is required to attain full correction of the spherical aberration. Hence the pick-off mirror (POM) is steerable in WFPC2, and three of the fold mirrors contain tip-tilt actuators. * CCD Technology: Many properties of WF/PC-1 and WFPC2 CCDs are compared in Table 4.1 on page 65. Many differences derive from the fact that the WF/PC-1 CCDs were thinned, backside illuminated devices whereas the WFPC2 CCDs are thick, frontside illuminated devices. WF/PC-1 CCDs were thinned and backside illuminated. The active silicon layer was a free-standing membrane somewhat less than 10 micro meter thick, with photons impinging directly on the silicon layer, without attenuation in the polysilicon gate structure built on the other ('front') side of the device. * Quantum Efficiency Hysteresis (QEH): The WF/PC-1 CCD's required a UV flood procedure and continuous cold temperatures to avoid QEH and hence non-linearity. A UV flood was performed early in the WF/PC-1 mission, but could not be repeated due to problems with the HST magnetometers. This in turn limited the temperature range allowable during decontaminations, since high temperatures would remove the UV flood, which in turn severely limited UV science capabilities. Some QE instability was also seen, particularly in the B band, due to changes in the UV flood. WFPC2 CCDs support multi-pinned phase (MPP) operation which eliminates quantum efficiency hysteresis. * Charge Transfer Efficiency: WF/PC-1 devices suffered from significant charge transfer efficiency (CTE) errors at image intensities below ~200 electrons per pixel. This was overcome by preflashing virtually all science images. WFPC2 devices have very little CTE error, and hence no preflash is used. Low-level charge traps are present in the WFPC2 devices, and result in a loss of about 4 percent of the signal when a star image is clocked down through all rows of the CCD. In the presence of background, the effect is reduced. For most applications, CTE is negligible or calibratable and pre-flash exposures are not required. * Detector MTF: The WFPC2 Loral devices do suffer from poorer CCD detector MTF than the WF/PC-1 CCDS, perhaps caused by scattering in the frontside electrode structure. The effect is to blur images and decrease the limiting magnitude by about 0.5 magnitudes. * Flat field quality: WFPC-1 CCDs were chemically thinned devices and therefore varied in thickness across the field-of-view causing large features in the flat fields. WFPC2 CCDs are un-thinned and the intrinsic response is uniform to ~3 percent across the field. * DQE: The WFPC2 CCDs have intrinsically lower QE than WF/PC-1 CCDs above 4800A, which is due to attenuation by frontside electrode structures. * Gain switch: WF/PC-1 had only a single analog-to-digital converter gain setting of 8 e- DN^-1 which saturated at about 30,000e-. Two gains are available with WFPC2: a 7 e- DN^-1 channel which gives reasonable sampling of the 5e- read noise, and which saturates at about 27,000e-, and a 14 e- DN^-1 channel which saturates at about 53,000e- and extends the useful dynamic range. * Quantization: The systematic analog-to-digital converter errors that were present in the low order bits on WF/PC-1 have been largely eliminated, contributing to a lower effective read noise in WFPC2. * Calibration Channel: WF/PC-1 contained a solar UV flood channel which was physically in the location of the present WFPC2 calibration channel. This transmitted solar UV light into the camera to provide a UV flood capability. * Entry Port: The WF/PC-1 camera was sealed by an afocal MgF2 window immediately behind the shutter. The WFPC2 entry port is open. * Chronographic Capability: WF/PC-1 contained a low reflectance spot on the pyramid (known as the Baum spot) which could be used to occult bright objects. This has been eliminated from WFPC2, since the spherical aberration severely reduces its utility. * Contamination Control: Since launch, WF/PC-1 suffered from the accumulation of molecular contaminants on the cold (-87 degrees C) CCD windows. This molecular accumulation resulted in the loss of FUV (1150-2000A) throughput and attenuation at wavelengths as long as 5000A. Another feature of the contamination was the "measles" - multiple isolated patches of low volatility contamination on the CCD window. Measles were present even after decontamination cycles, when most of the accumulated molecular contaminants were boiled off by wanning the CCDS. In addition to preventing UV imaging, these molecular contamination layers scattered light and seriously impacted the calibration of the instrument. WFPC2 has far less contamination than WF/PC-1 owing to pre-launch cleaning and bake-out procedures, careful design of venting paths to protect the optical bench area, and inclusion of Zeolite molecular absorbers in the design. There is now a decrease in throughput of about 30 percent per month at 1700A, but monthly decontamination procedures completely remove this material. This throughput drop is also highly predictable and can be calibrated out during photometric analyses. 1.5 Organization of this Handbook A description of the instrument is contained in Chapter 2. The filter set Is described in Chapter 3. CCD performance is discussed in Chapter 4. A description of the Point Spread Function is given in Chapter 5. The details necessary to estimate exposure times are described in Chapter 6. A summary of observation strategies is given in Chapter 7. Data products and standard calibration methods, and calibration plans are summarized in Chapter 8. A complete list of references is given in Chapter 9. This document summarizes the performance of the WFPC2 as known in April 1996 after two years of on-orbit calibration. Observers are encouraged to contact the STScI Help Desk, the WFPC2 WWW pages (see section 1.7 "WFPC2 Handbook on the WWW" below), and STScI Instrument Scientists for the latest information. 1.6 What's New in Version 4.0 Major revisions since Version 3.0 may be summarized as follows: * Comparisons to FOC, STIS, and NICMOS. * Observation Strategies: A new chapter (Ch. 7) has been added specifically to assist observers in preparing Phase II proposals. * Exposure Time Estimation: Ch. 6 has been largely re-written. Signal-to-noise ratio equations for Poisson-limited, background-limited, and generalized cases are now included for point sources (both with PSF fitting and aperture photometry) and extended sources. The WWW on-line Exposure Time Calculator program is briefly described. A new Appendix gives representative SNR values for various exposures of stellar, power law, and emission line sources. * CCD Performance: Material on dark current and CTE (charge transfer efficiency) has been updated. * Calibration: Material on UV throughput, dark current calibration, flat fielding, and impact of focus variations on photometry has been updated. Cycle 6 calibration proposals are described. * Other changes include addition of an index and an acronym list. 1.7 WFPC2 Handbook on the WWW This Handbook will appear on the WFPC2 pages accessible at: http://www.stsci.edu/ftp/instrument_news/WFPC2/top_wfpc2.html and will be updated as new information becomes available. 1.8 The Help Desk at STScI STScI maintains a Help Desk whose staff quickly provide answers to any HST-related topic, including questions about WFPC2 and the Cycle 7 proposal process. The Help Desk staff has access to all of the resources available at the Institute. They maintain a database of frequently asked questions and answers, so that many questions can be answered immediately. The Help Desk staff can also provide copies of STScI documentation, in either hardcopy or electronic form, including Instrument Science Reports and Instrument Handbooks. Questions sent to the Help Desk during normal business hours are usually answered within one hour. Questions received outside normal business hours will be answered within the first two hours of the next business day. Usually, the Help Desk staff will reply with the answer to a question, but occasionally they will need more time to investigate the answer. In these cases, they will reply with an estimate of the time needed to reply with the full answer. We ask that you please send all initial inquiries to the Help Desk. If your question requires a WFPC2 Instrument Scientist to answer it, the Help Desk staff will put a WFPC2 Instrument Scientist in contact with you. By sending your request to the Help Desk, you are guaranteed that someone will provide a timely response. To contact the Help Desk at STScI: * Send e-mail: help@stsci.edu * Phone: 1-410-338-1082 The Space Telescope European Coordinating Facility (ST-ECF) also maintains a Help Desk. European users should generally contact the ST-ECF for help; all other users should contact STScI. To contact the ST-ECF Help Desk in Europe: * Sende-mail: stdesk@eso.org. 1.9 Further Information The material contained in this Handbook is derived from ground tests and design information obtained by the IDT and the engineering team at JPL, and from on-orbit measurements. Other sources of information are listed below. A complete list of references appears on page 197. * HST Phase II Proposal Instructions, (Version 8.0, 15 December 1995). This document may be requested by e-mail from help@stsci.edu. * HST Data Handbook, (Version 2.0, December 1996). This document may be requested by e-mail from help@stsci.edu. * Calibrating Hubble Space Telescope: Post Service Mission (I 995). These document may be requested by e-mail from help@stsci.edu. * STSDAS Calibration Guide, (November 1991). This document may be requested by e-mail from help@stsci.edu. * The Reduction of WFIPC Camera Images, Lauer, T., P.A.S.P. 101, 445 (1989). * The Imaging Performance of the Hubble Space Telescope, Burrows, C. J., et. al., Ap. J. Lett., 369, L21 (1991). * Interface Control Document (ICD) 19, "PODPS to STSDAS" * Interface Control Document (ICD) 47, "PODPS to CDBS" * The Wide Field/Planetary Camera in The Space Telescope Observatory, J. Westphal and the WF/PC-1 IDT, IAU 18th General Assembly, Patras, NASA CP-2244 (1982). * The WFPC2 Science Calibration Report, Pre-launch Version 1.2, J. Trauger, editor, (1993). [IDT calibration report] * White Paper for WFPC2 Far-Ultraviolet Science, J. T. Clarke and the WFPC2 IDT (1992). This document may be requested by e-mail from help@stsci.edu. * The Performance and Calibration of WFPC2 on the Hubble Space Telescope, Holtzman, J., et al., P.A.S.P., 107, 156 (1995). * The Photometric Performance and Calibration of WFPC2, Holtzman, J., et al., P.A.S.P., 107, 1065 (1995). * The Institute's WFPC2 World Wide Web page at address: http://www.stsci.edu/ftp/instrument_news/WFPC2/wfpc2_top.html * The Institute's WFPC2 Space Telescope Analysis Newsletter (STAN), which is distributed monthly via e-mail, and provides notification of any changes in the instrument or its calibration. To subscribe, send e-mail to help@stsci.edu. ------------------------------------------------------------------------------ CHAPTER 2: Instrument Description In This Chapter... Science Objectives WFPC2 Configuration, Field-of-View, and Resolution Overall Instrument Description Quantum Efficiency Shutte Serial Clocks Overhead Times CCD Orientation and Readout Calibration Channel 2.1 Science Objectives The scientific objective of the WFPC2 is to provide photometrically and geometrically accurate images of astronomical objects over a relatively wide field-of-view (FOV), with high angular resolution across a broad range of wavelengths. WFPC2 meets or exceeds the photometric performance of WF/PC-1 in most areas. The goal is 1 percent rms photometric accuracy, which means that the relative response in all 8OO x8 OO pixels per CCD must be known to better than 1 percent through each filter, and that standard calibrations be done at this level. Currently, the absolute calibration in the primary broadband photometric filters is accurate at around the 2 percent level, and is expected to continue to improve. Success in this area is dependent on the stability of all elements in the optical train, particularly the CCDs and filters. The narrow point spread function is essential to all science programs being conducted with the WFPC2, because it allows one to both go deeper than ground based imagery, and to resolve smaller scale structure with higher reliability and dynamic range. Further, many of the scientific goals which originally justified the HST require that these high quality images be obtained across a wide field-of-view. The Cepheid distance scale program, for example, cannot be accomplished without a relatively wide field-of-view. A unique capability of the WFPC2 is that it provides a sustained, high resolution, wide field imaging capability in the vacuum ultraviolet. Considerable effort has been expended to assure that this capability is maintained. Broad passband far-UV filters, including a Sodium Wood's filter, are included. The Wood's filter has superb red blocking characteristics. Photometry at wavelengths short of 3000A is improved through the control of internal molecular contamination sources and the ability to put the CCDs through warm-up decontamination cycles without loss of prior calibrations. While the WFPC2 CCDs have lower V-band quantum efficiency than the WF/PC-1 chips, for many applications this is more than made up for by the lower read noise, and by the intrinsically uniform flat field. For example, these characteristics are expected to increase the accuracy of stellar photometry, which was compromised by uncertainty in the flat field in WF/PC-1. 2.2 WFPC2 Configuration, Field-of-View, and Resolution The field-of-view and angular resolution of the wide field and planetary camera is split up as follows (see Section 4.2 for more details on CCDS): Table 2.1: Summary of Camera Format. Camera Pixel and CCD Field-of-View Pixel Scale F/ratio Format ------------------------------------------------------------------------------ Wide Field 800 x 800 2.5' x 2.5' ~100 Milli- 12.9 x 3 CCDs (L-shaped) arcseconds Planetary 800 x 800 35" x 35" ~46 milli- 28.3 x 1 CCD arcseconds ------------------------------------------------------------------------------ Figure 2.1: Wide Field Planetary Camera Concept Illustration. The calibration channel, and pick-off mirror mechanisms are not shown. 2.3 Overall Instrument Description The Wide-Field and Planetary Camera, illustrated in Figure 2.1, occupies the only radial bay allocated to a scientific instrument. Its field-of-view is centered on the optical axis of the telescope and it therefore receives the highest quality images. The three Wide-Field Cameras (WFC) at F/12.9 provide an "L" shaped field-of-view of 2.5 x 2.5 arcminutes with each 15 gm detector pixel subtending 0.10" on the sky. In the Planetary Camera (PC) at F/28.3, the field-of-view is 35" x 35", and each pixel subtends 0.046". The three WFCs undersample the point spread function of the Optical Telescope Assembly (OTA) by a factor of 4 at 5800A in order to provide an adequate field-of-view for studying galaxies, clusters of galaxies, etc. The PC resolution is over two times higher. Its field-of-view is adequate to provide full-disk images of all the planets except Jupiter (which is 47" in maximum diameter). The PC has numerous extra-solar applications, including studies of galactic and extra-galactic objects in which both high angular resolution and excellent sensitivity are needed. In addition to functioning as the prime instrument, the WFPC2 can be used for target acquisition in support of other HST instruments, or for parallel observations. Figure 2.2 shows the optical arrangement (not to scale) of the WFPC2. The central portion of the OTA F/24 beam is intercepted by a steerable pick-off mirror attached to the WFPC2, and is diverted through an open entry port into the instrument. The beam then passes through a shutter and filters. A total of 48 spectral elements and polarizers are contained in an assembly of 12 filter wheels. Then the light falls onto a shallow-angle, four-faceted pyramid located at the aberrated OTA focus, each face of the pyramid being a concave spherical surface. The pyramid divides the OTA image of the sky into four parts. After leaving the pyramid, each quarter of the full field-of-view is relayed by an optical flat to a Cassegrain relay that forms a second field image on a charge-coupled device (CCD) of 8OO x 8OO pixels. Each detector is housed in a cell that is sealed by a MgF2 window. This window is figured to serve as a field flattener. The aberrated HST wavefront is corrected by introducing an equal but opposite error in each of the four Cassegrain relays. An image of the HST primary mirror is formed on the secondary mirrors in the Cassegrain relays. (The fold mirror in the PC channel has a small curvature to ensure this, and is why the PC magnification changed from F/30 in WF/PC-1 to F/28.3 in WFPC2.) The spherical aberration from the telescope's primary mirror is corrected on these secondary mirrors, which are extremely aspheric. The point spread function is then close to that originally expected for WF/PC-1. Figure 2.2: WFPC2 Optical Configuration The single most critical and challenging technical aspect of applying a correction is assuring exact alignment of the WFPC2 pupils with the pupil of the HST. If the image of the HST primary does not align exactly with the repeater secondary, the aberrations no longer cancel, leading to a wavefront error and comatic images. An error of only 2 percent of the pupil diameter produces a wavefront error of 1/6 wave, leading to degraded spatial resolution and a loss of about 1 magnitude in sensitivity to faint point sources. This error corresponds to mechanical tolerances of only a few microns in the tip/tilt motion of the pick-off mirror, the pyramid, and the fold mirrors. Mechanisms inside WFPC2 allow optical alignment on-orbit; these are necessary to insure correction of the OTA spherical aberration. The beam alignment is set with a combination of the steerable pick-off mirror and actuated fold mirrors in cameras PC1, WF3 and WF4. The 47 degree pick-off mirror has two-axis tilt capabilities provided by stepper motors and flexure linkages, to compensate for uncertainties in our knowledge of HST's latch positions (i.e., instrument tilt with respect to the HST optical axis). These latch uncertainties would be insignificant in an unaberrated telescope, but must be compensated for in a corrective optical system. In addition, three of the four fold mirrors, internal to the WFPC2 optical bench, have limited two-axis tilt motions provided by electrostrictive ceramic actuators and invar flexure mountings. Fold mirrors for the PC1, WF3, and WF4 cameras are articulated, while the WF2 fold mirror has a fixed invar mounting. A combination of the pick-off mirror and fold mirror actuators has allowed us to correct for pupil image misalignments in all four cameras. Since the initial alignment, stability has been such that mirror adjustments have not been necessary. The mechanisms are not available for GO commanding. The WFPC2 pyramid cannot be focused or rotated. WFPC2 is focused by moving the OTA secondary mirror, and then COSTAR (or any future science instruments) is adjusted to achieve a common focus for all the HST instruments. The four CCDs provide a 1600 x 1600 pixel field-format; three of the 800 x 800 CCDs have 0.1" pixels (WFC), and one has 0.046" pixels (PC). The CCDs are physically oriented and clocked so that the pixel read-out direction is rotated approximately 90 degrees in succession (see Figure 1.1 on page 2). The (1,1) pixel of each CCD array is thereby located near the apex of the pyramid. As a registration aid in assembling the four frames into a single picture, a light can be turned on at the pyramid to form a series of eleven fixed artificial "stars" (known as Kelsall spots or K-spots) along the boundaries of each of the quadrants. This calibration is normally done in a separate exposure. The K-spot images are aberrated and similar in appearance to the uncorrected HST PSF. The relative alignment of the four channels has been more accurately determined from star fields, and is stable over long periods, but the K-Spot images are useful for verifying the stability. Figure 2.3: Cooled Sensor Assembly Each CCD is a thick frontside-illuminated silicon sensor, fabricated by Loral Aerospace. A CCD, mounted on its header, is hermetically packaged in a ceramic-tube body that is filled with 1.1 atmosphere of argon to prevent degradation of the UV sensitive phosphor, and sealed with the MgF2 field flattener. This complete cell is connected with compliant silver straps to the cold junction of a thermoelectric cooler (TEC). The hot junction of the TEC is connected to the radial bay external radiator by an ammonia heat pipe. This sensor-head assembly is shown in Figure 2.3. During operation, each TEC cools its sensor package to suppress dark current in the CCD. 2.4 Quantum Efficiency The WFPC2 provides useful sensitivity from 1150A to IIOOOK in each detector. The overall spectral response of the system is shown in Figure 2.4 (not including filter transmissions). The curves represent the probability that a photon that enters the 2.4m diameter HST aperture at a field position near the center of one of the detectors will pass all the aperture obscurations, reflect from all the mirrors, and eventually be detected as an electron in the CCD. The throughput of the system combined with each filter is tabulated in Table 6.1 and also shown in the Appendix. Figure 2.4: WFPC2 + OTA System Throughput. These measurements made on orbit are much more accurate than the pre-launch estimates, and are used consistently throughout this Handbook. The visible and red sensitivity of the WFPC2 is a property of the silicon from which the CCDs are fabricated. To achieve good ultraviolet response, each CCD is coated with a thin film of Lumogen, a phosphor. Lumogen converts photons with wavelengths less than 4800A into visible photons with wavelengths between 5100A and 5800,k, which the CCD detects with good sensitivity. Beyond 4800,k, the Lumogen becomes transparent and acts to some degree as an anti-reflection coating. Thus, the full wavelength response is determined by the MgF2 field flattener cutoff on the short-wavelength end and the silicon band-gap in the infrared at 1.1eV (~11OOOA). With the WFPC2 CCD sensors, images may be obtained in any spectral region defined by the chosen filter with high photometric quality, wide dynamic range, and excellent spatial resolution. The bright end of the dynamic range is limited by the 0.11 seconds minimum exposure time, and by the saturation level of the analog-to-digital converter (ADC) at the chosen gain, which is roughly 53000 (gain=14, though called ADT-GAIN=15 in RPS2) or 270OOe- (gain=7) per pixel. The maximum signal-to-noise ratio corresponding to a fully exposed pixel will be about 230. The faint end of the dynamic range is limited by photon noise, instrument read noise and, for the wide-band visible and infra-red filters, the sky background. Table 2.2 gives characteristic values of the expected dynamic range in visual magnitudes for point sources. The minimum brightness is given for an integrated S/N ratio of 3, and the maximum corresponds to CCD ADC saturation (selected as 530OOe-). The quoted values assume an effective bandwidth of 1000A at about 5600A (filter F569W). The planets and many other resolved sources are observable in this filter with short exposures even if their integrated brightness exceeds the 8.5 magnitude limit. Table 2.2: WFPC2 Dynamic Range in a Single Exposure. Configuration Exposure (seconds) Min. V Magnitude Max. V Magnitude ------------------------------------------------------------------------------ Wide Field 0.11 8.82 17.83 Wide Field 3000. 19.91 28.14 Planetary 0.11 8.40 17.47 Planetary 3000. 19.49 28.20 ------------------------------------------------------------------------------ 2.5 Shutter The shutter is a two-blade mechanism used to control the duration of the exposure. A listing of the possible exposure times is contained in Table 2.3. These are the only exposure times which can be commanded. Current policy is to round down non-valid exposure times to the next valid value. An exposure time of less than 0.11 seconds will therefore only result in a bias frame being taken. Some exposures should be split into two (CR-SPLIT) in order to allow cosmic ray events to be removed in post-processing. By default, exposures of more than 10 minutes are CR-SPLIT. If an exposure is CR-SPLIT, the exposure time is divided into two fractions and then rounded down. Normally the fractional split is 50 percent/50 percent but, unless constrained by the user with CR-TOLERANCE, the ratio may be up to 70 percent/30 percent, as allowed by the default CR-TOLERANCE=0.2. Note that some exposure times in the table do not correspond to commendable values when halved. In preparing a proposal containing an exposure that is to be CR-SPLIT, the simplest procedure to use in order to be sure of a given total exposure time, is to enter double a legal value, and impose CR-TOLERANCE = O. For the shortest exposure times, it is possible to reconstruct the actual time of flight of the shutter blades. Encoder disks, attached to the shutter blade arms, are timed by means of a photo-transistor. The maximum error is 5 milliseconds. The necessary information is contained in the WFPC2 engineering data stream, however, this information is not in the processed science header. Diffraction effects from the edges of the shutter blades affect the point spread function for very short exposures. It is advisable to use exposure times greater than 0.2 second when obtaining point spread functions in support of long exposure observations (see the WF/PC-1 IDT OV/SV Report, Chapter 9, for further discussion in the spherically aberrated case). The control of the initial opening of the WFPC2 shutter during an observation is held by the internal WFPC2 microprocessor in all cases. However, control over when the shutter is closed is held by the microprocessor only for exposures less than 180 seconds in duration. For longer exposures, control passes to the Application Processor (AP-17) in the NSSC-1 spacecraft computer. The consequence of this arrangement is that loss of guide star lock will result in the WFPC2 shutter being closed only for those observations with planned durations of 180 seconds or longer. The AP-17 always controls the shutter closing if the serial clocks are enabled during the exposure (CLOCKS=YES), which then has a minimum planned duration of 1 second, and exposures are rounded to the nearest second. If guide star lock is reacquired prior to the end of the planned observation time, the shutter will reopen to obtain a portion of the planned integration. As discussed in the next section, CLOCKS=YES should generally not be used with exposures shorter than 30 sec., if 1 percent or better photometric accuracy is needed. Table 2.3: Quantized Exposure Times (Seconds). Exposure times where the PSF is affected by shutter flight time are underlined. Exposures normally without loss of lock checking are in italics. Times that are CR-Split by default are in boldface. Exposures that take more than one orbit even when CR-split are not normally accessible to GOs and are crossed (and exposures longer than 5400 seconds must be CR-Split). Exposure times that should not be used when CLOCKS=YES are shaded. 2.6 Serial Clocks The serial transfer registers of the CCDs can be kept running during an exposure (CLOCKS=YES), or run only during the readout (CLOCKS=NO, the default). The serial clocks are sometimes used on very bright targets where extensive blooming of the up and down the CCD columns is expected. CLOCKS=YES causes charge which blooms to the ends of the CCD to be read out and disposed of, thus preventing it from flowing back into the image. They will be useful when any single CCD column contains in excess of ~10^8 electrons. Note that the serial clocks do not actually suppress the blooming process, instead they merely remove any charge that blooms to the top or bottom of the CCD. For most circumstances, we recommend CLOCKS=NO. The reasons for this recommendation are: 1. CLOCKS=YES is not widely used, so anomalies may exist or develop that we are not aware of Also, this mode is not as well calibrated as CLOCKS=NO (although we expect the calibration to be independent of the state of the clocks). 2. The shutter open time when CLOCKS=YES can have significant errors. In this mode, there are delays of up to 0.25 seconds in opening the shutter (which are not present when CLOCKS=NO). This means that for exposures of less than ~30 seconds, there may be photometric errors greater that 1 percent, unless special precautions are taken in the data reduction. Furthermore, if a non-integral exposure time is specified in the proposal, it will be rounded to the nearest second. See below for details. On the other hand: 1. We do advise CLOCKS=YES if you expect star images to be so saturated that a significant amount of charge will bleed off the chip during the exposure. This would mean that you expect much more than 10^8 electrons from at least one star in the exposure (more than 1000 pixels would be saturated). Otherwise the charge can be detected in other parts of the image. 2. One advantage of CLOCKS=YES is that the overhead time is 1 minute less for exposures longer than 180 seconds. This can be significant if you have a large number of exposure times in the 3 to 10 minute range. 3. Unlike the original WF/PC-1, we do not see a significant variation of WFPC2 dark level with CLOCKS=YES. In summary: 0 - 0.8 sec CLOCKS=NO is required. 1 - 30 sec Use CLOCKS=NO (or attempt photometric corrections during the analysis of the data). 20 - 180 sec Use CLOCKS=NO unless more than 108 detected electrons from a single star are expected. 180+ sec Use CLOCKS=NO unless more than 108 detected electrons are expected, or if you need to save 1 minute of overhead. While exposure times are corrupted for CLOCKS=YES, and are not accurately reported in the image headers, correct values can be computed. Details are as follows: 1. Non-integer exposure times <3 minutes are rounded to the nearest integer (e.g., 1.2 see and 1.4 see will actually be 1.0 see long, 3.5 see exposures take 4.0 see). This roundoff is due to the way the spacecraft computer monitors the take-data flag (AP-17 uses its own integer-based timecode). This rounding is reflected properly in the header keywords (keywords UEXPODUR, EXPSTART, EXPEND, EXPTIME, and EXPFLAG in the .c0h file headers, or UEXPODUR and CMD_EXP in the .shh headers). 2. All CLOCKS=YES exposures are also shortened by either 0.125 or 0.250 seconds. This decrease in exposure time is not reflected in the file headers; the amount depends upon which shutter blade was in place at the start of the exposure. The decrease in exposure time is due to the manner in which the application processor (AP-17) in the spacecraft computer operates the shutter blades. When CLOCKS=NO (default), the WFPC2 microprocessor opens the shutter, the AP-17 closes the shutter, and the exposure time is as requested. However, with CLOCKS=YES, the AP-17 opens the shutter, first blade A, then blade B. If blade A is closed at the start of the exposure, the actual exposure begins 0.125 seconds after the AP-17 issues the blade command. If blade B is closed at the exposure start, the exposure starts 0.250 seconds later (after the AP-17 sends the open-A command followed by open-B). The shutter in place at exposure start is given in the SHUTTER keyword in the .coh file. 2.7O verhead Times Efficient use of the WFPC2 requires an understanding of the overhead times of the instrument. In this section, the various causes of overhead are presented in a manner that should allow the user to make a fairly accurate prediction of the cost in time of a given sequence of exposures. This information is provided for completeness and background. Guidelines in the Phase I proposal instructions and RPS2 should be followed to develop Phase I and II proposals, respectively. (See also "Choosing Exposure Times" on page 152.) 1. Telescope alignments. A telescope alignment is, in practice, a set of images uninterrupted by a POS-TARG target position change or the end of orbit. The start of an alignment requires 1 minute overhead in order to synchronize timing with a major frame (all commands to the instrument take place on major frames which last 1 minute). The end of alignment uses one minute for tape recorder overhead. If scans are being performed, another minute of overhead is required and, if images are requested in real-time, another 2 minutes must be added to the alignment end. There are additional overheads at the start of each target visibility period associated with guide star acquisition (9 minutes), or reacquisition (6 minutes). 2. Filter changes. A filter change requires at least 1 minute, the use of 2 filters requires 2 minutes of overhead. Furthermore, since the filter history is lost across telescope alignments, at least one minute is spent on selecting the filter at an alignment start, regardless of the filter in place before the alignment. 3. CCD clearing. Clearing the CCD is done before every exposure and requires 14 seconds. This time is part of the first major frame of the exposure. Therefore, time taken for a given exposure (excluding all other overheads) is the exposure time plus 14 seconds rounded up to the next integral minute. For example, all legal exposure times up to 40 seconds inclusive cost one minute. 4. CCD readout. The readout time for an exposure is one minute. An additional minute is required for exposures 180 sec. or longer, taken with CLOCKS=NO. This extra minute can be saved by using CLOCKS=YES, but this is not generally recommended (see section 2.6, "Serial Clocks", on page 27). If the exposure is CR-SPLIT, the readout overheads (calculated with the split exposure times) are of course doubled. There is normally no overhead time advantage in reading out a subset of the CCDS. The exception is if the WFPC2 readout occurs in parallel with the operation of a second instrument, when at least 2 minutes is required to read all 4 CCDS. 5. Spatial scans/Dithering. Spatial scans are specially designed sequences of images taken with telescope pointing changes periodically placed between successive images. The paintings in a spatial scan must either be equally spaced points on a single line, or a grid of points formed by the intersection of two sets of equally spaced parallel (but not necessarily mutually perpendicular) lines. Scans avoid the much larger alignment overheads associated with the use of POS-TARG special requirements. Dithering is the use of small spatial displacements to allow better removal of chip defects and/or the reconstruction of sub-pixel resolution. During Phase II the user will be given access to canned" dithering routines which will allow him/her to avoid many of the tricky details involved in planning spatial scans. The overheads in these canned routines is the same as that of a user-planned spatial scan. 6. The overhead of a spatial scan is similar to that of a sequence of images taken in one alignment; however, at least one minute of overhead is required for each change in pointing. Furthermore, an extra minute of overhead is incurred at the end of the scan and typically about 1 minute of overhead is used at the beginning of the scan positioning the first image. In summary, it is not possible to schedule exposures in different filters less than 3 minutes apart: commands to the WFPC2 are processed at spacecraft "major frame" intervals of one minute. A filter wheel may be returned to its "clear" position and another filter selected in one minute. An exposure takes a minimum of one minute, and a readout of the CCDs takes one or two minutes depending on the exposure time. Hence a simple exposure requires a minimum of 3 minutes. Table 2.4: Instrument Overheads. The first and last exposures of an alignment contain extra overheads. Overhead Type Time (min.) Overhead ------------------------------------------------------------------------------ First exposure 1 Major frame uncertainty, clock synchronization First exposure 1 To put in initial filter Per image: 1 Per filter change Per image: int(t)+1 t=(shutter-open time in seconds +14 seconds)/60 Per image: 1 If CLOCKS=NO (default) and exposure >= 180 sec Per image: 1 Readout Per image: 1 If image done in parallel with another instrument Last exposure 1 Tape overhead Last exposure 2 If data requested down in real-time Last exposure 1 If a scan was done ------------------------------------------------------------------------------ 2.8 CCD Orientation and Readout The relation between the rows and columns for the four CCDs is shown in Figure 1.1 on page 2. Each CCD's axes are defined by a 90 degree rotation from the adjacent CCD. If a 4-CCD image is taken and then each subimage is displayed with rows in the "X" direction and columns in the "Y" direction, each successive display would appear rotated by 90 degree from its predecessor. Table 2.5: Inner Field Edges. The CCD X,Y (Column, Row) numbers given vary at the 1-2 pixel level because of bending and tilting of the field edge in detector coordinates due to the camera geometric distortions. Start Vignetted Field 50 percent Start Unvignetted Camera (Zero Illumination) Illumination Field (100 percent Illumination) ------------------------------------------------------------------------------ PC1 X>O and Y>8 X>44 and Y>52 X>88 and Y>96 WF2 X>26 and Y>6 X>46 and Y>26 X>66 and Y>46 WF3 X> IO and Y>27 X>30 and Y>47 X>50 and Y>67 WF4 X>23 and Y>24 X>43 and Y>44 X>63 and Y>64 ------------------------------------------------------------------------------ Figure 1.1 on page 2 also illustrates the projected orientation of the WFPC2 CCDs onto the sky. The beam is split between the four cameras by a pyramid-shaped mirror in the aberrated HST focal plane. In an effort to insure images from the four CCDs can be reassembled into a single image without gaps, there is a small overlap region on the sky between each CCD and its neighbors (see also Figure 3.11 on page 62). On the CCDs this region appears as a blank 4 1 shadow" region along the X~0 and Y~0 edges of each CCD; the exact limits of this region are given in Table 2.5 for each CCD. Because the OTA beam is aberrated at the pyramid mirror, the edges of the shadow region are not sharp, but instead there is a gradual transition from zero to full illumination on each CCD. The width of this vignetted region is essentially that of the aberrated OTA beam (~5"). Table 2.5 gives approximate limits of this vignetted region on each CCD. Note that astronomical sources in the vignetted region are imaged onto two or more CCDS. The WFPC2 has two readout formats, namely full single pixel resolution (FULL Mode), and 2x2 pixel summation (AREA Mode which is obtained by specifying the optional parameter SUM=2x2 as described in the Proposal Instructions). Each line of science data is started with two words of engineering data, followed by 800 (FULL) or 400 (AREA) 16-bit positive numbers as read from the CCDs (with 12 significant bits). In FULL Mode the CCD pixels are followed by 11 "bias" words ("over-clocked" pixels), yielding a total of 813 words per line for 800 lines. In AREA Mode, there are 14 bias words giving a total of 416 words per line for 400 lines. Either pixel format may be used to read out the WFC or PC. These outputs are reformatted into the science image and extracted engineering data files during processing in the HST ground system prior to delivery to the observer. The advantage of the AREA Mode (2x2) on-chip pixel summation is that readout noise is maintained at 5 e- RMS for the summed (i.e., larger) pixels. In addition, more images will fit onto the spacecraft tape recorder. This pixel summation is useful for some photometric observations of extended sources particularly in the UV. Note, however, that cosmic ray removal is more difficult in AREA Mode. The readout direction along the columns of each CCD is indicated by the small arrows near the center of each camera field in Figure 1.1 on page 2 (see also Figure 3.11 on page 62). Columns and rows are parallel and orthogonal to the arrow, respectively. Each CCD is read out from the comer nearest the center of the diagram, with column (pixel) and row (line) numbers increasing from the diagram center. In a saturated exposure, blooming will occur almost exclusively along the columns because of the MPP operating mode of the CCDS. Diffraction spikes caused by the Optical Telescope Assembly and by the internal Cassegrain optics of the WFPC2 are at 45 degrees to the edges of the CCDS. Unless specified otherwise in the Phase 2 proposal, the default pointing position when all 4 CCDs are used is on WF3, approximately 10" along each axis from the origin. Observations which require only the field-of-view of a single CCD are best made with the target placed near the center of a single CCD rather than near the center of the 4 CCD mosaic. This results in a marginally better point spread function, and avoids photometric, astrometric, and cosmetic problems in the vicinity of the target caused by the overlap of the cameras. Even so, for such observations the default operational mode is to read out all four CCDS. This policy has resulted in serendipitous discoveries, and sometimes the recovery of useful observations despite pointing or coordinate errors. On the other hand, any combination of 1, 2 or 3 CCDs may be read out in numerical order (as specified in the Proposal Instructions). This partial readout capability is not generally available to GOs, although it can be used if data volume constraints mandate it, after discussion with the WFPC2 instrument scientists. It does not result in a decrease in the readout overhead time but does conserve limited space on the HST on-board science tape recorder. The capacity of the present science tape recorder is slightly over 7 full (4-CCD) WFPC2 observations and 18 single CCD WFPC2 observations on a single tape recorder side (of two). Switching sides of the tape recorder without a pause will result in the loss of part of a single CCD readout. Since an interval of about 30 minutes must normally be allowed for the tape recorder to be copied to the ground, readout of only a subset of the WFPC2 CCDS, or use of AREA mode, can be advantageous when many frames need to be obtained in rapid succession. Note, however, that the Solid State Recorder planned for installation during the 1997 Service Mission is capable of hold well over one hundred 4-CCD WFPC2 images. This capability should be phased-in during Cycle 7, and will lead to relaxation of the above data rate restrictions for Cycle 7 proposals. Multiple exposures may be obtained with or without interleaved spacecraft repointings and filter changes without reading the CCDs (READ=NO). These would then be followed by a readout (READ=YES). Note that WFPC2 must be read out at least once per orbit. 2.9 Calibration Channel An internal flat field system provides reference flat field images over the spectral range of WFPC2. These are provided by a "calibration channel" optical system mounted outside the main shroud of WFPC2. The system consists of a series of lamps and diffusers, and a flip mirror which directs the beam into the WFPC2 entrance aperture. The lamp set contains Tungsten incandescent lamps with spectrum shaping glass filters and a Deuterium UV lamp. The flat field illumination pattern is fairly uniform for wavelengths beyond about 1600A. Short of 1600A the flat field is distorted due to refractive MgF2 Optics- In practice, the flat fields used routinely to calibrate WFPC2 data have been generated by combining flats taken with an external stimulus in thermal vacuum testing with Earth "streak" (unpainted) flats to give the low frequency terms in the OTA illumination pattern. The calibration channel is used primarily to check for internal instrumental stability. ------------------------------------------------------------------------------ CHAPTER 3: Optical Filters In This Chapter... Introduction Choice of Broad Band Filters Linear Ramp Filters Redshifted [OII] Quad Filters Polarizer Quad Filter 50 Methane Quad Filter 53 Wood's Filters Red Leaks in UV Filters Apertures 3.1 Introduction A set of 48 filters are included in WFPC2 with the following features: 1. It approximately replicates the WF/PC-1 "UBVRI" photometry series. 2. The broad-band filter series extends into the far UV. 3. There is a Stromgren series. 4. A Wood's filter is available for far-UV imaging without a red leak. 5. There is a 1 percent bandpass linear ramp filter series covering 3700-9800A. 6. The narrow-band series is uniformly specified and well calibrated. The filters are mounted in the Selectable Optical Filter Assembly (SOFA) between the shutter and the reflecting pyramid. The SOFA contains 12 filter wheels, each of which has 4 filters and a clear "home" position. A listing of all simple optical elements in the SOFA mechanism, and the location of each element (by wheel number 1-12, and position 1-4) is given in Table 3.1. Wheel number 1 is located closest to the shutter. The categories of simple filters (F) are long-pass (LP), wide (W), medium (M), and narrow (N). Most of these filters are either flat single substrates or sandwiches. The filter complement includes two solar blind Wood's filters, F160AW and F16OBW. F16OBW is used in all science observations because the other filter has some large pinholes that lead to significant red leak. In addition to the above complement of broad and narrow-band filters, WFPC2 features a set of three specialized quadrant (quad or Q) filters in which each quadrant corresponds to a facet of the pyramid, and therefore to a distinct camera relay. There is one quad containing four narrow-band, redshifted [OII] filters with central wavelengths from 3763A to 3986A, one quad with four polarizing elements (POL) with polarization angles, 0, 45, 90 and 135 degrees, and one quad with four methane (CH4) band filters with central wavelengths from 5433A to 8929A. The polarizer quad filter can be crossed with any other filter over the wavelength range from 2800A to 8000A, with the exception of the Methane Quad and Redshifted [OII] Quad which share the same wheel. The SOFA also contains four linearly variable narrow-band ramp (FR) filters (in the twelfth wheel - closest to the focus). The quad and ramp filters are listed in Table 3.2. In Table 3.1 and Table 3.2, each of the type "A" filters is equivalent to inserting 5 mm of quartz in terms of optical path length, with compensation for wavelength such that focus is maintained on the CCDS. A configuration with no filters in the beam results in out-of-focus images and generally will not be used. With the exception of the quad polarizer and blocking (Type "B") filters, all filters are designed to be used alone. Type "B" filters introduce no focus shift, so they can be used in combination with any type "A" filter. All combinations where the number of type "A" filters is not unity will result in out-of-focus images. The image blur resulting from two or zero type "A" filters at visible wavelengths is equivalent to 2.3 mm defocus in the F/24 beam, which corresponds to 1/5 wave RMS of defocus at 6328A, and a geometrical image blur of 0.34". While this is a large defocus, the images are still of very high quality compared to seeing limited images. Some such combinations may be scientifically attractive. For example, the Wood's filter may be crossed with another UV filter to provide a solar blind passband (although the efficiency will be low). The mean wavelength, lambda bar, is similar to that defined in Schneider, Gunn and Hoessel (Ap. J. 264, 337). The width is the FWHM of a Gaussian filter with the same second moment, and is reasonably close to the FWHM. The values tabulated here do not include the CCD DQE or the transmission of the OTA or WFPC2 optics (as given in Figure 2.4). In Chapter 6, the corresponding quantities are given including the effect of the other optical elements and the CCD DQE. Figure 3.1 summarizes the normalized transmission curves for the simple filters and narrow-band quad filters. It does not include curves for the polarizing quad, or the linear ramp filters which are documented in sections 3.2 and 3.4 respectively. Individual filter transmission curves are shown in the Appendix starting on page 201. Figure 3.1 divides the filters into the following groups: 1. Long pass filters designed to be used in combination with another filter. 2. Wide bandpass filters with FWHM -25 percent of the central wavelength. 3. Approximations to the UBVRI sequence, generally with wider bandpasses, designed for use on faint sources. 4. A photometric set of approximations to UBVRI passbands (see Haffis et al. 1991, A.J. 101, 677). Note, however, that the WFPC2 UBVRI series is not the Johnson-Cousins photometric series, neither is it identical with the WF/PC-1 series. See Chapter 6 for detailed comparisons. Table 3.1: WFPC2 Simple Filter Set. The effective wavelength, width, and transmission quoted are defined precisely in Chapter 6, but here are quoted without the system (OTA-WFPC2) response. T = Type W = Wheel S = Slot IN = In WF/PC-1? lb(A) = lambda-bar Angstroms dlb(A) = delta lambda-bar Angstroms P-T = Peak T percent P-l = Peak lambda Angstroms Name T W S Notes IN lb(A) dlb(A) P-T P-l ------------------------------------------------------------------------------ F122M A 1 4 H Ly alpha-Red Leak Y 1292 263.5 18.9 1240 F13OLP B 2 1 CaF2 Blocker (zero focus) N 3847 5568.1 94.7 4344 F160AW A 1 3 Woods A - redleak from pinholes N 1471 457.2 10.1 1403 F16OBW A 1 2 Woods B N 1471 457.2 10.1 1403 F165LP B 2 2 Suprasil Blocker (zero focus) N 4327 5505.7 92.1 3109 F170W A 8 1 - N 1689 434.9 30.3 1667 F185W A 8 2 - N 1907 302.9 23.7 1849 F218W A 8 3 Interstellar feature N 2136 355.9 21.3 2091 F255W A 8 4 - N 2557 408.2 15.5 2483 F300W A 9 4 Wide U N 2924 727.6 51.8 2760 F336W A 3 1 WFPC2 U, Stromgren u Y 3327 370.7 80.3 3447 F343N A 5 1 Ne V N 3430 24.3 19.7 3433 F375N A 5 2 [OII] 3727 RS Y 3736 26.2 17.2 3736 F38OW A 9 1 - N 3934 694.7 65.6 3981 F39ON A 5 3 CN N 3889 45.3 37.8 3886 F410M A 3 2 Stromgren v N 4088 146.9 70.1 4098 F437N A 5 4 [OIII] Y 4369 25.2 52.0 4368 F439W A 4 4 WFPC2 B Y 4292 464.4 67.3 4176 F450W A 10 4 Wide B N 4445 925.0 92.4 5061 F467M A 3 3 Stromgren b N 4682 171.5 75.3 4728 F469N A 6 1 He II Y 4695 24.9 52.4 4699 F487N A 6 2 H beta Y 4865 25.8 58.6 4863 F502N A 6 3 [OIII] Y 5012 26.8 63.8 5009 F547M A 3 4 Stromgren y (but wider) Y 5454 486.6 91.3 5361 F555W A 9 2 WFPC2 V Y 5252 1222.5 94.8 5151 F569W A 4 2 F555W generally preferred^a Y 5554 965.8 94.2 5309 F588N A 6 4 He I & Na I (NaD) Y 5892 49.1 91.5 5895 F606W A 10 2 Wide V Y 5843 1578.7 98.3 6183 F622W A 9 3 - Y 6157 935.4 95.6 6034 F631N A 7 1 [OI] Y 6306 30.8 85.7 6302 F656N A 7 2 H alpha Y 6562 22.0 77.9 6561 F658N A 7 3 [NII] Y 6590 28.5 79.7 6592 F673N A 7 4 [SII] Y 6733 47.2 87.0 6733 F675W A 4 3 WFPC2 R Y 6735 889.4 97.9 6796 F702W A 10 3 Wide R Y 6997 1480.7 98.5 6539 F785LP A 2 3 F814W generally preferred^a Y 9366 2094.7 96.0 9960 F791W A 4 1 F814W generally preferred^a Y 8006 1304.2 99.7 8081 F814W A 10 1 WFPC2I Y 8269 1758.0 98.4 8386 F850LP A 2 4 - Y 9703 1669.5 95.5 10026 F953N A 1 1 [SIII] N 9546 52.5 95.6 9528 F1042M A 11 2 - Y 10443 610.9 95.2 10139 ------------------------------------------------------------------------------ a. Filters F555W and F814W are generally preferred, as they are part of the "standard" WFPC2 filter set, and will tend to have slightly better photometric calibration. See "Choice of Broad Band Filters" on page 40. Table 3.2: WFPC2 Quad and Ramp Filters. Segments of the UV and CH4 quads are labeled here by their usual physical designations (A, B, C, and D); see following sections for filter and aperture names which are to be used in writing Phase 11 proposal. The quad polarizer is represented for both parallel and perpendicular polarization to its polarization direction, which is different in each quadrant. T = Type W = Wheel S = Slot IN = In WF/PC-1? lb(A) = lambda-bar Angstroms dlb(A) = delta lambda-bar Angstroms P-T = Peak T percent P-l = Peak lambda Angstroms Name T W S Notes IN lb(A) dlb(A) P-T P-l ------------------------------------------------------------------------------ FQUVN-A A 11 3 Redshifted [OII] 375 N 3767 73.4 24.0 3769 FQUVN-B A 11 3 Redshifted [OII] 383 N 3831 57.4 30.5 3827 FQUVN-C A 11 3 Redshifted [OII] 391 N 3915 59.3 38.9 3908 FQUVN-D A 11 3 Redshifted [OII] 399 N 3993 63.6 43.4 3990 FQCH4N-A A 11 4 CH4 543 N 5433 38.0 83.8 5442 FQCH4N-B A 11 4 CH4 619 N 6193 44.0 83.7 6202 FQCH4N-C A 11 4 CH4 727 N 7274 51.0 89.7 7279 FQCH4N-D A 11 4 CH4 892 N 8929 64.0 91.2 8930 POLQ-par B 11 1 Pol angle N 5427 5797.6 90.8 10992 0,45,90,135 degrees POLQ-per B 11 1 Pol angle N 7922 6653.4 89.7 10992 0,45,90,135 degrees FR418N A 12 1 3700-4720 N W W/75 ~20-50 W FR533N A 12 2 4720-6022 N W W/75 ~40-50 W FR68ON A 12 3 6022-7683 N W W/75 ~60-80 W FR868N A 12 4 7683-9802 N W W/75 ~70-85 W ------------------------------------------------------------------------------ 5. Medium bandpass filters with FWHM ~10 percent of the central wavelength, including an approximation to the Stromgren photometric series. 6. Narrow bandpass filters for isolating individual spectral lines or bands. 7. Redshifted [OII] and CH4 narrow bandpass quad filters. Note that the UV filters have some degree of "red leak," which is quantified in Chapter 6 where the system response is included. A passband calibration is maintained in the calibration database system (CDBS). It has been updated following on orbit calibrations. The ground based calibration of the narrow-band filters' central wavelengths has not been corrected for temperature effects and is therefore accurate to about 2A. Because of this, it is not advisable to place narrow emission lines at the half power points of such filters and expect to predict the throughput to high accuracy. The standalone software package XCAL, or SYNPHOT running under IRAF, can be used to access these calibrations which are available on the Institute's WWW page. Figure 3.1: Summary of Normalized Filter Curves. 3.2 Choice of Broad Band Filters A number of different choices are possible on WFPC2 in order to approximate the Johnson-Cousins system often used in ground based observing. These choices differ in throughput, wavelength fidelity, color transformability, and cosmetics The HST science program as a whole benefits if a standard set can be agreed upon by the community for broad band photometry. This will allow theoretical isochrones and other models to be published in the standard system, and allow ready comparison of the results from different observers. Furthermore, although all filters will be calibrated photometrically and with flat fields, a core set must be chosen for monitoring the instrument both photometrically and in imaging performance. There was a substantial consensus between the accepted Cycle 4 GO programs and the WF/PC-1 and WFPC2 science teams that F336W, F439W F555W, F675W, and F814W should be the preferred set to approximate the Johnson Cousins U, B, V, R, I passbands. These filters fon-n the basis for the WFPC2 broad band photometric system. As will be seen from the figures in Section 8.8 (page 172), the preferred set is accurately transformable with the exception of the U bandpass. On the other hand, there are situations where concerns such as maximum throughput must override the above arguments. For example, filters F300W, F450W, F606W, and F814W were chosen for the Hubble Deep Field (HDF), due to their wider bandpasses. 3.3 Linear Ramp Filters The linear ramp filters are designed for naffow-band absorption and emission line imaging of moderately extended objects. Each filter is divided into four parallel strips where the central wavelength across each strip varies by approximately 6 percent. Each CCD pixel is mapped to a unique central wavelength with a FWHM bandwidth of approximately 1.3 percent of the central wavelength. The maximum size of an object which can be imaged at a given wavelength is approximately 13" and is determined by the width of the strips and the image size at the filter. The cumulative wavelength range of the four linear ramp filters is 3710A to 9762A. Originally intended for a four WFC configuration, the linear ramp filters require partial rotation of the SOFA wheels to +15, -18 and -33 degrees from their nominal positions, to recover wavelength regions which would fall on the PC camera or otherwise be lost. There will be vignetting at some wavelengths for these partial rotations. 3.3.1 Spectral Response A JPL Memorandum (DFM #2031, 1992) gives the results of a prediction scheme to locate and quantify the passbands of the four WFPC2 ramp filters, FR418N, FR533N, FR68ON and FR866N. The results are summarized here. Laboratory (room temperature) measurements of the passbands of the four ramp filters were made at five equally spaced intervals on each of the four ramp stripes on each filter for a total of 80 passband measurements. The laboratory measurements were made with a narrow beam and were then integrated over an annular area of the filter to simulate the beam profile. The radius of the beam is 3.7 mm, or 13". The integration was carried out by assuming the nominal linear shift in wavelength with position, and that no significant changes in the passband shape occur across the beam. The integration makes the shape of the passband quite symmetrical. The resulting spectral response can be fitted to within a few percent with a Munson function: T = T_o /{1 + (1-a)x^2 + a(1-b)X^4 + ab(1-c)x^6 + abcx^8 where a, b and c are shape parameters, and O<=(a,b,c)<=I1; T_o is the peak transmission of the passband, T=T_o at x=O; x is related to wavelength X by x=(?,-XO)IH, T=TOI2 at x--1 (so H is the half width at half maximum). The parameters, (lambda_o, T_o, H, a, b, c) were then fitted to polynomial functions of position Y (which starts at 0 inches at the lower wavelength edge of each strip) to predict the filter response for areas of the filters between the tested points. Good quadratic fits are available for all the parameters except for To which requires a cubic. The results are given in Tables 3.3 through 3.6, which give the polynomial fit coefficients for the ramp filter parameters. The table entries, except for the first line, are used as parameter = A_o + A_1Y + A_2Y_2 + A_3Y_3. The short wavelength side of the filter is opposite for alternate ramps. The first line in each table gives the Y position as a function of X. If the polynomial fit predicts a, b, or c < 0 or > 1 then the quantities are set to 0 or 1 respectively. Use of these fits should be restricted to objects near the center of the ramp, otherwise the beam will combine light from adjacent ramps. The fit should also not be used within 13" of the end of the ramp. There is enough wavelength overlap between ramps that the extreme ends need not be used, except at the very lowest and highest wavelengths. Figure 3.2 shows the fit parameter To as a function of lambda_o for all 16 ramp filter strips. Figure 3.3 shows 2H/lambda_o. Figure 3.2: Ramp Filter Peak Transmission. The four line types correspond to the four different filters (each containing four ramps). Table 3.3: Ramp Filter FR418N Parameters. Quantity A0 A1 A2 A3 ------------------------------------------------------------------------------ Ramp I Position -26.1083 .00713888 .0000 Wavelength 3657.7 138.7 .6178 Peak transmission -.01667 .2188 .04138 -.03489 Half width at half max 21.95 -.8347 2.143 a .2120 .002857 .002596 b 1.181 -.8138 .3535 c .3301 -.3715 .3825 Ramp 2 Position -24.2554 .00625704 .0000 Wavelength 3876.9 158.6 .5472 Peak transmission .1660 .2288 -.1080 .004005 Half width at half max 21.50 3.315 -.7079 a .1592 -.003687 -.0008497 b .7938 .2355 -.09124 c .9306 .01366 .007458 Ramp 3 Position -24.7145 .00598254 .0000 Wavelength 4130.5 168.8 -.7389 Peak transmission .1352 .6200 -.5226 .1529 Half width at half max 22.09 1.306 -.1181 a .2300 .05586 -.03044 b 1.096 -.3185 .1396 c 1.276 -1.279 .5721 Ramp 4 Position -23.4440 .00536340 .0000 Wavelength 4371.3 185.8 .2913 Peak transmission .3189 .1287 -.01160 -.001712 Half width at half max 25.62 1.015 .1161 a .3123 -.2055 .09535 b .9222 .1167 -.04673 c 1.033 -.1356 .05660 ------------------------------------------------------------------------------ Table 3.4: Ramp Filter FR533N Parameters. Quantity A0 A1 A2 A3 ------------------------------------------------------------------------------ Ramp 1 Position -26.7670 .00572115 .0000 Wavelength 4677.7 177.3 -1.125 Peak transmission .5450 -.3612 .3623 -.1281 Half width at half max 25.67 .3168 .8873 a -.009839 .4644 -.2039 b .31511 .9473 -.4516 c -.3379 2.788 -1.46 Ramp 2 Position -24.6600 .00498393 .0000 Wavelength 4948.4 199.2 .6484 Peak transmission .4546 .4188 -.5456 0.1548 Half width at half max 32.10 -1.204 3.171 a .1678 -.02726 .09521 b .9345 .1935 -.1224 c .9571 .02919 -.009393 Ramp 3 Position -24.5038 .00465985 .0000 Wavelength 5257.3 217.9 -1.481 Peak transmission .4944 -.1714 .1890 -0.0631 Half width at half max 34.03 5.078 -1.347 a .3851 -.06264 .003163 b .5605 .6642 -.2751 c .9665 .05543 -.03654 Ramp 4 Position -25.5182 .00455886 .0000 Wavelength 5596.9 220.9 -.6938 Peak transmission .5058 -.2715 .3203 -.1230 Half width at half max 35.06 -2.856 2.382 a -.06553 .2253 -.08275 b 1.043 -.1190 .02889 c 1.162 -.4910 .2059 ------------------------------------------------------------------------------ Table 3.5: Ramp Filter FR68ON Parameters. Quantity A0 A1 A2 A3 ------------------------------------------------------------------------------ Ramp 1 Position -21.8962 .00370137 .0000 Wavelength 5916.0 269.4 .3460 Peak transmission .1198 1.005 -.4015 -.00162 Half width at half max 41.50 -5.873 4.038 a .1743 -.05050 .06481 b .8320 .3326 -.1858 c .9682 -.09110 .05122 Ramp 2 Position -22.6919 .00360750 .0000 Wavelength 6290.8 275.6 .7184 Peak transmission .7918 -.02034 .1086 -.05945 Half width at half max 39.48 2.120 .3703 a .05596 .3034 -.1333 b 1.017 -.27026 .04560 c .7244 .8326 -.5107 Ramp 3 Position -22.0719 .00330755 .0000 Wavelength 6673.5 301.6 .3321 Peak transmission .9494 -1.008 1.161 -.3777 Half width at half max 42.81 .8193 .4269 a .1038 .09020 -.02747 b .8415 .3045 -.1930 c 1.017 -.1732 .07463 Ramp 4 Position -24.7447 .00346462 .0000 Wavelength 7141.9 289.3 -.2999 Peak transmission .4823 .4479 -.07484 -.05868 Half width at half max 44.72 .8952 -.0756 a .1612 -.01167 .01355 b .2708 1.077 -.4757 c .9941 -.02694 .01685 ------------------------------------------------------------------------------ Table 3.6: Ramp Filter FR868N Parameters. Quantity A0 A1 A2 A3 ------------------------------------------------------------------------------ Ramp 1 Position -23.2685 .00308029 .0000 Wavelength 7555.5 320.4 1.906 Peak transmission .7524 -.3328 .4543 -.1343 Half width at half max 49.32 1.742 .4914 a .2958 -.3877 .2465 b 1.321 -.9156 .3666 c .3762 1.668 -.9499 Ramp 2 Position -22.9766 .00286673 .0000 Wavelength 8014.3 350.5 -.7500 Peak transmission .8204 -.3368 .3815 -.1057 Half width at half max 54.17 1.579 .2196 a .05832 .7525 -.3625 b .4582 .8433 -.4350 c .6422 .3247 -.193 Ramp 3 Position -22.6085 .00265657 .0000 Wavelength 8510.7 375.6 .3706 Peak transmission .5817 -.1920 .4517 -.1627 Half width at half max 55.19 -.7459 1.433 a .5422 -.2444 .03545 b 1.420 -1.176 .4814 c .4257 -.2522 .1777 Ramp 4 Position -23.2142 .00256976 Wavelength 9034.3 387.2 .8722 Peak transmission .6241 .2403 -.1230 .02829 Half width at half max 59.69 2.167 -.1996 a .2376 -.01879 -.00864 b .9670 .02456 -.00477 c .7829 .03750 .02393 ------------------------------------------------------------------------------ 3.3.2 Target Locations In Figures 3.4 and 3.5 we show the correspondence between central wavelength and location in the focal plane for the nominal and rotated filter positions. The selection of filter and aperture for the linear ramp filters is transparent to the user who is required only to specify the linear ramp filter name LRF and a central wavelength. Each central wavelength is assigned to a unique filter and CCD location. Following on-orbit testing of WFPC2 a revised table of linear ramp filter wavelengths has been compiled and is shown in Table 3.7. For each wavelength listed, there is a minimum 10" diameter unvignetted field-of-view. Some wavelengths can be obtained with several different settings of the ramps, however, for simplicity, the redundant wavelengths have been removed from the table. Note that this table supports observation with the PC and a new +15 degree rotation of the filter wheel. Table 3.8 lists wavelengths which are available, but with some compromise in data quality, so as to avoid gaps in wavelength coverage. Most of these wavelengths are observed slightly off the central wavelength of the passband. This implies a slightly reduced throughput (see estimates of the light reduction in the table), and some additional difficulties in flattening the data to remove variations in the passband across the target. A few other wavelengths are observed slightly off the unvignetted centerline of the ramps, and these are indicated by note "FOV" in Table 3.8. Again, this vignetting will present some additional complications when calibrating the data. Further details regarding the ramp filter wavelengths and apertures will be made available in a separate instrument science report. We note that an interactive tool is available on the WFPC2 WWW pages which will compute target locations for LRF observations. The user inputs either the central wavelength or the target location in the field-of-view, and the other quantity is returned. 3.3.3 LRF Photometric Calibration As of this writing, the preferred method of flat fielding LRF data is to use a narrow band flat observed nearby in wavelength. This will remove pixel-to-pixel effects, as well as effects of distortion and vignetting in the cameras, while avoiding the complications of pinholes on the LRFs and spurious variations due to the spectrum of the flat field light source. Conversion of counts to source flux is best achieved by using the SYNPHOT synthetic photometry package. An LRF filter setting is simply specified by including "LRF#xxxx" in the OBSMODE, where xxxx is the central wavelength specified on the Phase II proposal. Figure 3.4: FR418N and FR533N Wavelength Mapping. Figure 3.5: FR68ON and FR868N Wavelength Mapping. Table 3.7: Aperture Locations and Wavelengths for Ramp Filters. Start (A) End (A) Filter CCD/ x1 y1 x2 y2 Aperture (pix) (pix) (pix) (pix) ------------------------------------------------------------------------------ 3710 3800 FR418N WF4-FIX 750 736.8 161.5 737.7 3800 3878 FR418N33 WF3-FIX 669.5 559.2 395.1 128.9 3881 3907 FR418N18 PC1-FIX 402.3 225 515.4 579.5 3907 3929 FR418N33 WF2-FIX 128.4 286.7 250.1 209.9 3929 4008 FR418N18 WF2-FIX 562.7 233 130.1 367.1 4008 4038 FR418N PC1-FIX 541.3 632.7 543.3 256.5 4038 4100 FR418N18 WF3-FIX 425.3 130.8 532.4 469.9 4100 4177 FR418N WF4-FIX 309 276.2 750.3 275.5 4186 4210 FR418P15 WF4-FIX 596.5 515.9 469.4 482.1 4210 4308 FR418N WF3-FIX 248.2 665.9 252.7 128.5 4308 4337 FR418P15 PC1-FIX 690.2 264.6 598.4 599.9 4337 4446 FR418N WF2-FIX 127.9 247.6 725.4 255.7 4446 4550 FR418N WF2-FIX 691.7 716.2 180.6 709.2 4550 4571 FR418P15 WF2-FIX 230 253.8 130.7 225.8 4593 4720 FR418N WF3-FIX 713.7 125.6 708.5 749.9 4746 4863 FR533N WF3-FIX 689.3 748.9 694.4 135.5 4884 4900 FR533P15 WF2-FIX 128.3 205.1 209 227.9 4900 5013 FR533N WF2-FIX 153.6 689.6 745.9 697.7 5013 5020 FR533N18 WF2-FIX 693.4 642.4 662.9 651.8 5020 5153 FR533N WF2-FIX 737.3 236.6 130 228.4 5153 5176 FR533P15 PC1-FIX 637.9 614.9 698.6 393.3 5188 5310 FR533N WF3-FIX 233.5 127.4 228.8 684.7 5310 5335 FR533P15 WF4-FIX 482.8 505.5 593.1 534.9 5339 5450 FR533N WF4-FIX 750.9 294.7 277.2 295.5 5450 5528 FR533N1g WF3-FIX 504.4 445.3 404.1 127.6 5528 5566 FR533N PC1-FIX 585.3 277.5 583.4 632.3 5566 5671 FR533N18 WF2-FIX 124@ 1 348.8 552.3 216.1 5671 5700 FR533N33 WF2-FIX 224.8 203.2 122.3 267.7 5700 5741 FR533N18 PC1-FIX 558.8 577 444.9 220.1 5743 5910 FR533N33 WF3-FIX 370.8 126.5 745.9 714.9 5910 6007 FR533N WF4-FIX 333.8 747.6 738.8 746.9 6007 6192 FR680N WF2-FIX 750.3 706.9 122.9 698.4 6192 6208 FR680P15 WF2-FIX 177.1 228.4 124.9 213.6 6238 6409 FR680N WF3-FIX 703.6 128.1 698.8 708.2 6409 6584 FR680N WF3-FIX 237.8 705.6 242.6 127 6590 6631 FR680P15 PC1-FTX 699.1 315.3 620.9 601.2 6631 6800 FR680N WF2-FIX 125.9 237.5 684.5 245.1 6800 6921 FR680N18 WF2-FIX 480.1 248 129.9 356.6 6921 6976 FR680N PC1-FIX 563.3 639.2 565.3 274.6 6976 7061 FR680N18 WF3-FIX 413.2 126 490.8 371.7 7061 7241 FR680N WF4-FIX 203 286.4 748.3 285.6 7251 7420 FR680N WF4-FIX 749.6 743.5 213.3 744.3 7420 7600 FR680N33 WF3-FIX 688.9 608.4 381.6 126.4 7605 7658 FR680N18 PC1-FIX 427 230 538.9 580.6 7658 7690 FR680N33 WF2-FIX 126.2 276.1 212.1 222 7690 7830 FR868N WF4-FIX 711.5 751.3 316.5 751.9 7830 8072 FR868N33 WF3-FIX 728.2 705.8 360.9 129.7 8077 8140 FR868N18 PC1-FIX 471.5 231 589.7 601.5 8140 8300 FR86SN18 WF2-FIX 527.6 213.2 126.2 337.6 8300 8362 FR868N PC1-FIX 605.4 644.1 607.3 287.9 8362 8460 FR868N18 WF3-FIX 393.1 126.1 470.6 371.7 8460 8661 FR868N WF4-FIX 196.9 305.7 724.7 304.9 8661 8910 FR868N WF3-FIX 218.3 731.6 223.4 125.3 8945 8980 FR868P15 PC1-FIX 701.1 467.5 651.9 647.3 8980 9200 FR868N WF2-FtX 142.7 218.5 678.2 225.9 9200 9415 FR868N WF2-FIX 668.4 686.5 162.2 679.6 9415 9456 FR868P15 WF2-FIX 219.9 220.5 127 194.2 9501 9762 FR868N WF3-FIX 684.3 135.4 679.2 750.2 ------------------------------------------------------------------------------ Table 3.8: Vignetted Wavelengths for Ramp Filters. The right column gives the maximum throughput reduction (in percent) for these settings where the target must be placed away from the optimal location on the filter glass. "FOV" denotes settings where transmission is optimal, but the usable field-of-view is reduced below 10" to the indicated diameter (in arcseconds). Start (A) End (A) Filter CCD/ x1 y1 x2 y2 Max Aperture (pix) (pix) (pix) (pix) percent Light Loss ------------------------------------------------------------------------------ 3878 3881 FR418N18 PC1-FIX 402.3 225.0 402.3 225.0 2 4177 4182 FR418N WF4-FIX 750.3 275.5 750.3 275.5 3 4182 4186 FR418P15 WF4-FIX 596.5 515.9 596.5 515.9 2 4571 4582 FR418P15 WF2-FIX 130.7 225.8 130.7 225.8 13 4582 4593 FR418N WF3-FIX 713.7 125.6 713.7 125.6 13 4720 4733 FR418N WF3-FIX 708.5 749.9 708.5 749.9 14 4733 4746 FR533N WF3-FIX 689.3 748.9 689.3 748.9 14 4863 4873 FR533N WF3-FIX 694.4 135.5 694.4 135.5 8 4873 4884 FR533P15 WF2-FIX 128.3 205.1 128.3 205.1 8 5176 5183 FR533P15 PC1-FIX 698.6 393.3 698.6 325.9 FOV~9" 5183 5188 FR533N WF3-FIX 233.5 127.4 233.5 127.4 2 5335 5337 FR533P15 WF4-FIX 593.1 534.9 593.1 534.9 1 5337 5339 FR533N WF4-FIX 750.9 294.7 750.9 294.7 1 5741 5743 FR533N33 WF3-FIX 370.8 126.5 370.8 126.5 1 6208 6221 FR680P15 WF2-FIX 124.9 213.6 124,9 213.6 8 6221 6238 FR680N WF3-FIX 703.6 128.1 703.6 128.1 11 6584 6587 FR680N WF3-FIX 242.6 127.0 242.6 127.0 1 6587 6590 FR680P15 PC1-FIX 699.1 294.3 699.1 315.3 FOV~9" 7241 7246 FR680N WF4-FIX 748.3 285.6 748.3 285.6 2 7246 7251 FR680N WF4-FIX 749.6 743.5 749.6 743.5 2 7600 7602 FR680N33 WF3-FIX 381.6 126.4 381.6 126.4 1 7602 7605 FR680N18 PC1-FIX 427.0 230.0 427.0 230.0 1 8072 8074 FR868N33 WF3-FIX 360.9 129.7 360.9 129.7 1 8074 8077 FR868N18 PC1-FIX 471.5 231.0 471.5 231.0 1 8910 8920 FR868N WF3-FIX 223.4 125.3 223.4 125.3 2 8920 8945 FR868P15 PC1-FIX 701.1 339.1 701.1 467.5 FOV~7" 9456 9478 FR868P15 WF2-FIX 127.0 194.2 127.0 194.2 13 9478 9501 FR868N WF3-FIX 684.3 135.4 684.3 135.4 13 ------------------------------------------------------------------------------ 3.4 Redshifted [OII] Quad Filters The redshifted [OII] quad filter was designed to map onto a four-faceted WFC configuration. A partial SOFA wheel rotation of -33 degrees is required to move filter quadrant 1 (3763A) into WF2 and WF3, with some vignetting of both camera fields. The projections of the redshifted [OII] filter settings FQUVN and FQUVN33 onto the field-of-view are essentially identical to those of the POLQ and POLQN33 filters, respectively (Figure 3.6). The vignetted regions are similar, and the location of aperture FQUVN33 is identical to that of POLQN33. The nominal and rotated filter wheel positions for the redshifted [OII] quad filter are each associated with different filter names. This allows pipeline calibration and database retrievals to proceed smoothly. The filter names are summarized in Table 3.9. The required central wavelength is selected by filter name and aperture location. Filter element FQUVN (Filter Quad Ultra Violet Narrow) has three possible apertures, each of which is nominally centered in one of the three WF channels and associated with a unique central wavelength. The filter element FQUVN33 corresponds to a single central wavelength. In addition to the filter name and aperture, a central wavelength is also requested in the proposal instructions in order to provide a consistency check. Aperture names are discussed further in section 3.9, "Apertures", on page 60. Table 3.9: Redshifted [OII] Quad Filter Elements. Filter Aperture FOV Mean Effective Name Name Location Quad Wavelength Width Comments ------------------------------------------------------------------------------ FQUVN WF2 WF2 D 3992 64 Nominal filter wheel position FQUVN WF3 WF3 C 3913 59 Nominal filter wheel position FQUVN WF4 WF4 B 3830 57 Nominal filter wheel position FQUVN33 FQUVN33 WF2 A 3765 73 Filter rotated -33 degrees ------------------------------------------------------------------------------ 3.5 Polarizer Quad Filter The polarizer quads were also designed to map onto a four-faceted WFC configuration and, consequently, also require a partial filter rotation of -33 degrees to move the filter quadrant 1 (nominal polarization angle 135 degrees) into WFCs 2 and 3 with some vignetting of both camera fields. Several additional partial rotations have been added to allow observations with different polarization angles on the same CCD. The polarizer quad may be used in several ways: by observing the target with each camera, by observing the target with the same camera using different partial rotations of the polarizer quad, or by observing the target with the same camera using different roll angles of the spacecraft. The first method has the drawback that calibration is complicated by uncertainties in the relative photometric calibration between cameras, while the second method uses the same camera but has non-optimal polarization angles and limited fields of view. The third method may present scheduling difficulties due to constraints on the spacecraft roll angle, and the need to rotate undersampled images. (See Biretta and Sparks 1995, "WFPC2 Polarization Observations: Strategies, Apertures, and Calibration Plans," WFPC2 Instrument Science Report 95-01.) The polarizer is designed for problems where large polarization angles are observed, and is not suitable for problems requiring precision of order 3 percent or better. The required polarization angle is selected by filter name and aperture location. The transmission of the quad polarizer is shown in Figure 3.7. The polarizer is afocal and must therefore usually be used with another filter which will largely define the shape of the passband. Table 3.10: Polarizer Quad Filter. Polarization angle 0 degrees lies along +X direction in Figure 3.11 Filter Aperture FOV Polarization Comments Name Name Location Angle ------------------------------------------------------------------------------ POLQ PC1 PC1 135 degrees Nominal filter wheel position POLQ WF2 WF2 0 degrees Nominal filter wheel position POLQ WF3 WF3 45 degrees Nominal filter wheel position POLQ WF4 WF4 90 degrees Nominal filter wheel position POLQN33 POLQN33 WF2 102 degrees Filter wheel rotated -33 degrees POLQP15 POLQP15P PC 15 degrees Filter wheel rotated +15 degrees POLQP15 POLQP15W WF2 15 degrees Filter wheel rotated +15 degrees POLQN18 POLQN18 WF2 117 degrees Filter wheel rotated -18 degrees ------------------------------------------------------------------------------ Figure 3.6: Polarizer Quads. The schematics show the filter projected onto the field-of-view for all rotated positions. Apertures are marked. Dashed lines indicate the central region of each quad which is free of vignetting and cross-talk. Greyscale images are VISFLATs of the polarizer with F555W. Figure 3.7: Polarizer Transmission for light polarized perpendicular (dotted curve) and parallel (solid curve) to the filter polarization direction. Methane Quad Filter The methane band quad filter, known as the jewel-quad, was designed for a four-faceted WF/PC configuration to permit imaging with both four WFC CCDs and four PC CCDS. WFC imaging is recovered for the first quadrant element of the filter (6193A) by a partial SOFA wheel rotation of -33 degrees which moves quadrant 1 into WF2 and WF3 with some vignetting of both camera fields. PC imaging with all four elements of the methane band jewel-quad cannot be recovered, but partial SOFA wheel rotations of -15 degrees and +15 degrees are implemented to recover two of the four methane band filters (8929A and 6193A). The +15 degree rotation of the filter wheel, however, results in some vignetting of PC1's field-of-view. The filter projections associated with the methane band jewel-quad are shown in Figure 3.8. Each of the four filter wheel positions are associated with unique filter names, as summarized in Table 3.11 on page 54. The required central wavelength is selected by filter name and aperture location. Filter element FQCH4N (Filter Quad Methane Narrow) has three possible apertures, each of which is located in one of the three WF channels and associated with a unique central wavelength, while FQCH4N33 is associated with one possible central wavelength. FQCH4N15 and FQCH4P15 are both associated with one central wavelength for PC1 observations. In addition to the filter name and aperture, a central wavelength is also requested in the proposal instructions in order to provide a consistency check. Table 3.11: Methane Band Quad Filter. The filter and aperture names should be specified on the Phase II proposal as shown here. Filter Aperture FOV Mean Effective Name Name Location Quad Wavelength Width (A) Comments (A) ------------------------------------------------------------------------------ FQCH4N FQCH4W2 WF2 A 5433 38 Nominal filter position FQCH4N FQCH4W3 WF3 D 8929 64 Nominal filter position FQCH4N FQCH4W4 WF4 C 7274 51 Nominal filter position FQCH4N33 FQCH4N33 WF2/WF3 B 6193 44 Filter rotated -33 degrees FQCH4N15 FQCH4NI5 PC1 B 6193 44 Filter rotated -15 degrees FQCH4P15 FQCH4PI5 PC1 D 8929 64 Filter rotated +15 degrees ------------------------------------------------------------------------------ 3.7 Wood's Filters WFPC2 features two solar-blind Wood's filters, for FUV (< 2000A) imaging. It was shown by Wood in the 1930s (Physical Optics, 1934, R. W. Wood) that thin layers of alkali metals transmit FUV wavelengths while providing very efficient long wavelength blocking due to the plasma frequency of the free electrons. Wood's filters have been built for WFPC2 at JPL using thin (5000A) layers of sodium sandwiched between two MgF2 substrates. These Wood's filters have a broad bandpass from 1200A to 2100A with visible-light transmission lower than 10-8. The best conventional UV filters exhibit visible-light transmission of 10-3 to 10-4. Many astronomical objects emit 10^4 to 10^7 visible photons for every FUV photon. In such cases, a Wood's filter (or "solar blind" detector as on STIS) is essential for FUV imaging so that the visible light leak does not dominate the observation. The main problem experienced with Wood's filters is long term instability. Sodium is a very reactive metal, and attempts to passivate the sodium layer have met with limited success. It is possible that as the Wood's filters age pinholes will form which transmit visible light. This transmitted light will appear as an increase in the background level at the focal plane. No indications of any degradation on-orbit have been observed. Figure 3.8: Methane Quad Filter. The mapping to the focal plane for nominal and rotated (-33, -15, and +15 degrees) SOFA positions is shown. Dashed lines indicate the limits of the unvignetted field-of-view on each quad. Figure 3.9: Wood's Filters. Greyscale flat field images show the field-of-view available with the two Wood's filter options F16OBW and F16OBN15. The Wood's filters can be used as a broadband filter, or in combination with the CaF2 long-pass filter to suppress geocoronal emission, or, crossed with one of the other UV filters, such as the suprasil blocker F165LP, to define a solar-blind UV photometric system. As discussed at the beginning of this chapter, the image will be out of focus in the last case. WFPC2's Wood's filters are circular with a clear aperture of 41 mm. Two similar Wood's filters (F160AW and F160BW) were mounted in SOFA wheel 1 to provide some redundancy. In Thermal Vacuum testing F160AW showed evidence for pinholes, which cause excessive red leak in some parts of its field. Therefore the preferred filter for far UV imaging with minimal red leak in WFPC2 is F160BW. In the nominal filter wheel position PC1 has a clear field-of-view but, there is significant vignetting in all three WFCS. A partial filter wheel rotation of -15 degrees produces a larger field-of-view in WF3, although some vignetting remains. The options are illustrated in Figure 3.9. The imaging performance of the Wood's filters is continually monitored for signs of aging such as visible light leaks. Additional partial rotations could be implemented in the future, to position an unaffected region of the filter into a WF or PCI, if necessary. The unvignetted filter projections associated with the two planned filter positions are shown schematically in Figure 3.9. Each filter position is associated with a unique name as summarized in Table 3.12. The filter name must be selected on the basis of whether a PC or WF3 observation is required. Table 3.12: Wood's Filters. The filter and aperture names should be specified on the Phase II proposal as shown below. Filter Aperture FOV Mean Effective Name Name Location Wavelength Width (A) Comments ------------------------------------------------------------------------------ F160BW PC1 PC1 1600 900 Nominal filter position F160BN15 F160BN15 WF3 1600 900 Filter rotated -15 degrees ------------------------------------------------------------------------------ 3.8 Red Leaks in UV Filters The "red leaks" in the UV filters are shown in Figure 3.10 for F122M F160BW (the new Wood's filter), F170W, F185W F218W F255W, F300W, and F336W. The presence of significant red leaks in the UV filters, together with the much greater sensitivity and wavelength coverage in the red part of the spectrum, makes calibration of UV observations difficult. Table 3.13 shows red leak estimates as a percentage of the total detected flux from de-reddened stellar sources, ordered by spectral type. In each column, the red leak is defined as the percentage of the detected flux longward of the cutoff wavelength in the second row. In the presence of interstellar reddening, the red leaks will be larger. Note that the SYNPHOT synthetic photometry package can be used to estimate the counts contributed by red leak for various particular situations, and for filters other than those plotted below. There is significant variation of the UV throughput due to build-up of molecular contaminants on the CCD windows, and decontamination procedures used to remove this contamination. See "Time Dependence of UV Response" on page 137. Figure 3.10: UV Filter Red Leaks. Includes the on-orbit measurements of system response. Table 3.13: Red Leak in UV Filters. A synthetic photometry calculation with de-reddened BPGS stellar spectra and system response from on-orbit data. 3.9 Apertures The WFPC2 camera configuration and filter set require a substantial number of apertures for full utilization of its capabilities. All possible aperture and filter combinations are given in Table 3.14. Each camera has an associated 'optimum' aperture close to the geometric center of its field-of-view (FOV). These positions have been adjusted to reflect CCD performance following SMOV and to allow for pyramid vignetting. The aperture designations are WF2, WF3, WF4, and PC1 for the individual cameras and WFFALL for the three-WFC combination. WFFALL is located close to the apex in WF3 (see Figure 3.11). Observers are expected to place small or unresolved targets on these apertures. Note that normally all four CCDs are read out even if a specific CCD is selected with an aperture. This is discussed in section 2.8, "CCD Orientation and Readout", on page 31. In cases where the observer does not want to use the current 'optimum' centers, a complimentary set of apertures has been implemented specifically for this purpose. These locations remain fixed and correspond roughly to the geometric center of each camera's field-of-view. They are designated WF2-FIX, WF3-FIX, WF4-FIX, PC1-FIX, and WFALL-FIX. Observers are expected to place extended targets on these apertures. An additional set of aperture names have been defined for use with the WFPC2 filters which require partial rotations. The characteristics and uses of these filters are discussed earlier in this Chapter. In the nominal filter position, the three WFC segments of the [OII], Methane and Polarizer quad filters can be selected with an aperture for each camera corresponding to the optimum or geometric camera centers. The partially rotated quad filters, which generally fall into more than one camera, have been assigned apertures in the camera which provides the largest clear aperture. The pixel coordinates of these apertures will be reviewed on a regular basis to reflect changes in CCD and filter cosmetics. There are no analogous fixed apertures for the partially rotated filter configurations. The aperture name is generally the same as the (rotated) filter name. For the Wood's filters, the nominal filter position is used for the PC1 FOV only, while the rotated filter position is used for WFC observations. The linear ramp filters are unique because the ultimate location of the target will be determined from the central wavelength specified and so an aperture name is not required. Occasionally the V2-V3 coordinates of the WFPC2 apertures are updated to correct slow drifts of the HST focal plane relative to the spacecraft (V1, V2, V3) system. Table 3.15 shows this history. The previous V2-V3 coordinates for any aperture can be derived by setting (V2_2,V3_2) to the values in Table 3.14, and then computing the earlier coordinates. Table 3.14: Aperture Definitions. The pixel coordinate system uses pixel numbers (row, column) for the CCD in use. See Figure 3.11: on page 62 or Figure 1.1: on page 2 for the definition of the V2-V3 coordinate system. CCD Pixel Coordinates^a Aperture Filter CCD Location X Y V2 V3 Name Name ------------------------------------------------------------------------------ PC1 PC Optimum center PC 420 424.5 2.160 -30.490 WF2 WF2 Optimum center WF2 423.5 414 -51.530 -5.920 WF3 WF3 Optimum center WF3 436^b 424.5b -0.150 48.470 WF4 WF4 Optimum center WF4 423 421 54.830 -6.320 WFALL WF3 Optimum near apex 133^c 149^c 2.020 7.920 PC1-FIX PC Fixed center PC 420 424.5 1.810 -30.900 WF2-FIX WF2 Fixed center WF2 423.5 414 -51.530 -5.920 WF3-FIX WF3 Fixed center WF3 416.5 424.5 1.230 47.100 WF4-FIX WF4 Fixed center WF4 423 421 54.830 -6.320 WFALL-FIX WF3 Fixed near apex 133^c 149c 2.020 7.920 FQUVN33 FQUVN33 WF2 Optimum for FOV 292 520 -49.924 10.802 P0LQN33 P0LQN33 WF2 Optimum for FOV 292 520 -49.924 10.802 P0LQN18 P0LQN18 WF2 Optimum for FOV 380^d 200^d -33.280 -17.717 P0LQP15P P0LQP15 PC Optimum for FOV 200 680 -13.057 -31.643 P0LQP15W P0LQP15 WF2 Optimum for FOV 500 260 -45.892 -22.069 FQCH4NW2 FQCH4N WF2 Optimum for FOV 602 608 -77.669 -5.102 FQCH4NW3 FQCH4N WF3 Optimum for FOV 602 608 0.928 72.995 FQCH4NW4 FQCH4N WF4 Optimum for FOV 640 386 67.687 11.323 FQCH4N33 FQCH4N33 WF2 Optimum for FOV 264 436 -41.997 6.950 FQCH4N15 FQCH4N15 PC Optimum for FOV 420 424.5 2.164 -30.494 FQCH4P15 FQCH4P15 PC Optimum for FOV 400 312 5.129 -26.221 F160BN15 F160BN15 WF3 Optimum for FOV 436 424.5 -0.153 48.470 ------------------------------------------------------------------------------ a. V2-V3 coordinates in effect after 1996 day 127 (May 6). CCD pixel positions unchanged. b. CCD pixel position in effect after 1994 day 101 (April 11). c. WFALL and WFALL-FIX "meta-chip" coordinates are (903,904). d. CCD pixel position in effect after 1995 day 86 (March 27). Table 3.15: Updates to (V2,V3) Positions of WFPC2 Apertures. ------------------------------------------------------------------------------ Date in Effect V2 V3 Rotation 1994 day 101 - 1996 day 105 V2_0 V3_0 PA_O 1996 day 105 - 1996 day 127 V2_1 = V3_1 = PA_1 = V2_0 - 0.12" V3_0 + 0.11" PA_0 + 0.14 degrees > 1996 day 127 V2_2 = V3_2 = PA_2 = V2_0 + 0.46" V3_0 + 0.39" PA_0 + 0.14 degrees ------------------------------------------------------------------------------ Figure 3.11: Precise CCD Alignments and Primary Aperture Locations. "FIX" apertures are in the same locations, unless otherwise indicated. Dashed lines show vignetted region along CCD boundaries. Short lines and "X"s in outer corners indicate directions of CCD bloom and OTA diffraction spikes, respectively Origin of the (V2, V3) system is at the origin of the plot axes, with V2, V3, and U3 exactly along diagonal lines as marked. (V2,V3) system is post-1996 day 127. CCDs have pixel (1,1) located where the four CCDs overlap. ------------------------------------------------------------------------------ CHAPTER 4: CCD Performance In This Chapter... Introduction Quantum Efficiency Dynamic Range Bright Object Artifacts Residual Image Quantum Efficiency Hysteresis Flat Field Response Dark Backgrounds Cosmic Rays Radiation Damage and Hot Pixels Charge Transfer Efficiency Read Noise and Gain Settings 4.1 Introduction The WFPC2 CCDs are thick, front-side illuminated devices, with a format of 8OOx8OO, 15x15 micro meter multi-pinned phase (MPP). MPP allows CCD exposure with the total inversion of all phases. The Si-SiO2 interface, at the surface of the CCD is pinned at the substrate potential, directing signal charge away from the Si-SiO2 interface states towards the buried n-channel. Figure 4.1 shows a schematic which illustrates the principle of MPP (modified from Janesick et al. 1989). The front-side Si-SiO2 interface significantly affects the performance of CCDS, so MPP operation yields many practical benefits including reduced dark noise, better charge transfer efficiency (CTE), rapid removal of residual images, excellent pixel-to-pixel uniformity, and improved radiation hardness. MPP technology has been demonstrated and characterized in both Loral (Janesick, et al., 1989) and Tektronix devices (Woodgate, et al., 1989). The CCD sensors for WFPC2 were made by Loral in 1991 and processed and packaged for flight at JPL. The Loral CCDs are illuminated from the 'front' surface, i.e., the light passes through the polysilicon gate structure overlying the 10 micro meter thick active silicon layer. Because the WFPC2 devices are front-side illuminated and supported by a bulk silicon substrate, the CCD surface is flat, which has reduced the uncertainties in the astrometric calibration to about the 1/10 pixel level. Figure 4.1: MPP Operating Principle. A schematic showing the ideal potential profile through a frontside illuminated CCD whose front surface is inverted with multi-pinned phase (MPP). The CCD consists of a polysilicon gate, which forms part of the electrode structure, a surface layer of oxidized silicon (SiO2) and the epitaxial layer which comprises p-doped silicon with an n-doped buried-channel for charge transfer. MPP pins the surface potential by populating the Si-SiO2 interface with holes. The holes passivate the Si-SiO2 interface states and create an electric field which directs signal charge away from the interface towards the buried n-channel. In this section the performance of the WFPC2 CCDs is reviewed, and compared to the WF/PC-1 devices. A summary of device characteristics is given in Table 4.1. 4.2 Quantum Efficiency The Loral CCDs are thick, front-side illuminated devices. This lowers their intrinsic QE, due to the absorption of incident light by the polysilicon electrode structure on the front-side surface of the CCD. We also note that due to its MPP operation, the QE of the Loral devices is stable without maintenance such as UV flooding. The front surfaces of the CCDs are overcoated with a Lumogen phosphor, which serves as the primary detection medium for photons shortward of about 4800A, down-converting these to 5100A - 5800A. Its long wavelength cutoff 4800A) is also well matched to a CCD's intrinsic sensitivity. The QE of the four flight WFPC2 CCDs is shown in Figure 4.2, which demonstrates the uniform UV response of 10-15 percent and a peak optical QE of 40 percent. This phosphor coating also produces an enhancement of DQE at visual wavelengths, since it acts as an anti-reflection coating. Table 4.1: Comparison of WF/PC-1 and WFPC2 CCDS. Parameter WF/PC-1^a WFPC2 ------------------------------------------------------------------------------ Device Ti Loral Architecture Thinned Thick Illumination back-side front-side Format 8OOx8OO 8OOx8OO Pixel size 15^2 gm 15^2 gm UV Phosphor Coronene Lumogen Dark rate 0.03 e- pixel^-1 s^-1(-87 C) ~0.0045 e- pixel^-1 s^-1(-88 C) Read noise 13e- RMS 5e- RMS Full well depth 40,000 e- ~90,000 e- Gain 8e- DN^-1 7e- DN^-1 or 14e- DN^-1 ADC range 12 bits (4096 DN) 12 bits (4096 DN) Full range (e-) ~30,000e- ~53,000e- QE 6000A 50 percent 35 percent QE 2500A 12 percent 15 percent WFC resolution 0.10" pixel^-1 0.0996" pixel^-1 PC resolution 0.043" pixel^-1 0.0455" pixel^-1 ------------------------------------------------------------------------------ a. WF/PC-1 data are available through the STScI data archives 4.3 Dynamic Range Linear full well capacity for these devices, clocked appropriately for the MPP mode, is approximately 90,000e- pixel^-1. Flight qualified ADCs with higher dynamic range (> 12 bits) were not available, so WFPC2 operates the two available ADCs at different gain factors, to take partial advantage of both the low read noise and large available full well depth. One channel has a gain of 14e- DN^-1, which significantly undersamples the CCD read noise (5 e- pixel^-1 RMS), and gives a digital full well of about 53,000e-. The other channel has a gain of 7e- DN^-1 which is comparable to the CCD read noise, and saturates at about 27,000e-. The choice of gain factor will be determined by the scientific objective. The 7 e- DN^-1 channel is best suited for faint object and UV imaging, where the lower CCD read noise will be most effective. For example, it should be used for UV imaging of planets or narrowband imaging of high redshift galaxies. The 14 e- DN^-1 channel has slightly higher effective read noise due to the quantization granularity, but can be used for programs where a signal level in excess of 27,000e- is required. Even when imaging faint sources, it may be desirable to retain the high signal-to-noise information on brighter field stars as a PSF reference. Use of the 14 e- DN^-1 channel also allows reasonable recovery of counts for isolated, saturated point sources by summing over the saturated pixels (assuming that the charge bleeding does not extend to the edges of the CCD). See Gilliland (1994). Figure 4.2: Pre-flight DQE Measurements on WFPC2 CCDS. The differences between the chips are probably due to systematic measurement error, and do not reflect a real difference in sensitivity. 4.4.1 Blooming Blooming up and down a CCD column occurs when more than about 90,000e- (the full well capacity) are collected in any pixel. When the pixel is full, the charge will flow into the next pixels along the column, and so on. The orientation of the bloomed column(s) on the sky depends on the readout direction of the particular CCD (see Figure 1.1 on page 2 or Figure 3.11 on page 62) and the roll angle of the spacecraft. This effect is visible in Figure 4.3 which shows a logarithmic stretch of the image resulting from a 100s exposure on a star of V magnitude 2.6 through filter F502N in the PC. Extreme overexposure of the Loral CCDs is not believed to cause any permanent effects, and therefore the WFPC2 does not have a bright object limit. The WFPC2 CCDs can be operated in a non-standard mode during the integration phase of an exposure, in order to limit the blooming to only those columns containing the bright sources. This is accomplished by operating the serial transfer register clocks during the integration (using the optional parameter CLOCKS as specified in the Proposal Instructions). See section 2.6, "Serial Clocks", on page 27 for details. 4.4.2 Horizontal Smearing During readout of a badly overexposed image, there is spurious charge detected by the readout electronics. The apparent brightness of the stellar halo is higher to the right of the saturated columns. This is particularly obvious at the bottom of the image in Figure 4.3 which is a region in the shadow of the pyramid edge. The horizontal "smearing" seen in highly saturated images can be modeled as an exponential function which decays over a few rows after a saturated pixel is encountered. The effect itself temporarily saturates after about ten saturated pixels (subsequent saturated pixels have no effect). The effect is twice as bad with gain 7 e- DN^-1 than with gain 14 e- DN^-1. This model only works on very highly saturated stellar images. In Figure 4.3, the image to the right side of the saturated columns is brighter than the left side; and the brightness increases as the number of saturated columns increases. This effect appears to be a signal which starts at a saturated pixel and decays over the next few rows, wrapping around as it does so. The signal is additive with each successive saturated pixel. Jumps are obvious when the number of saturated columns changes. The problem is a known characteristic of the amplifier electronics, and effort was made to minimize it during design. The increase in signal in rows with saturated pixels is also seen in the over-scan region (the over-scans are provided in ".xOd" files from the pipeline) An exponential function fits the effect reasonably well. An appropriate algorithm creates an array to contain the signal model. It searches through the uncalibrated image (with the over-scan region included) in the sequence in which the pixels are read out. When it encounters a saturated pixel, it adds an exponential function to the model array, beginning at that pixel. The function has the form s(x)=Ae^(-x/h), where x is the offset from the saturated pixel and only positive x values are included. The half-width, h, and amplitude, A, appear to vary from frame to frame and must be determined on the image itself. As more saturated columns are encountered in a row, the signal intensity builds up in the model image. The image can then be "improved" by subtracting the model from the raw image. The amplitude and half-width parameters can be obtained by trial and error. The typical parameters vary slightly for each chip. The amplitude per saturated pixel is typically 1.75 DN (gain 7) or 0.2 DN (gain 14). On the other hand the half-width at a gain of 14 is larger (h=1800) than at 7 (h=350). So the total integrated effect is about twice as bad at gain 7. A straightforward application of the above algorithm cleaned up most of the signal in rows which had a few saturated columns, but over-subtracted in rows with a large number. The algorithm can be modified to saturate by making the parameter A, which gives the peak contribution from a single saturated pixel, depend on the current level of the effect: A=A_0*(1-C/C_max). This implies that the correction is never larger than C., no matter how many saturated pixels are encountered. C. is approximately 14 DN for a gain of 7 and 10 DN for a gain of 14. The algorithm gives improvement only on highly saturated stellar images (where the star is saturated to 3 or 4 columns at the edges of the chip). On less saturated data, it over-subtracts significantly. This indicates that the problem is nonlinear, and therefore a general algorithm applicable to all data will be difficult to develop. 4.4.3 Diffraction Effects and Ghost Images Several other artifacts that are common in saturated stellar images are also obvious in Figure 4.3. The spider diffraction spikes caused by both the OTA spiders and internal WFPC2 spiders are at 45 degrees to the CCD columns in all cameras. The halo around the stellar image is well above the diffraction limit in intensity. Also there are ghost images which result from internal reflections in the filters and in the field-flatteners. These topics are discussed fully in the next Chapter. 4.5 Residual Image Residual images are seen in front-side-illuminated CCDs, and are associated with the front-side Si-SiO2 surface interface. When the full well is exceeded electrons can become trapped at the Si-SiO2 interface. This trapped charge is slowly released giving rise to residual images. Inverted phase operation (MPP) allows holes to recombine with the trapped electrons at the front-side interface, and so residual images dissipate in a matter of minutes. A second potential source of residual images, which occurs only in front-side-illuminated CCDS, is known as residual bulk image (RBI). Long wavelength photons can penetrate deeply enough to produce charge in the substrate. Most of this charge recombines rapidly (due to short carrier lifetimes), but some may diffuse into the epitaxial layer, where it can become trapped in epitaxial interface states. Residual images can occur as this charge is slowly released during an exposure. RBI is temperature sensitive since the bulk trapping time constants decrease with increasing temperature, The WFPC2 CCDs do exhibit RBI, but at -70 degrees C trapped charge rapidly escapes so that residual images disappear within 1000s. Driven by the WFPC2 electronics, the CCDs recover quickly from large over-exposures (100 times full well or more), showing no measurable residual images a half hour after the overexposure. For images exposed below the saturation level there is a very weak residual image due to charge trapping and charge transfer efficiency (CTE) problem. Measurements on 1800s dark frames interleaved with 2800s exposures of a star field yield a residual flux of 0.3 percent +- 0.1 percent of the original star flux, for stars with fluxes from 65 to 17,000 total counts. For typical star fields observed by WFPC2, these residual images are likely to be a problem only for stars that were saturated in a previous image, or for programs where long exposures in low throughput filters are taken immediately after highly exposes images. Hence, repeated exposures at the same CCD position should not lead to any appreciable systematic offset in photometry. CTE is further discussed in Section 4.1 1. Figure 4.3: Saturated Stellar Image Showing Horizontal Smearing. 4.6 Quantum Efficiency Hysteresis The problem of quantum efficiency hysteresis (QEH) due to back-side charge accumulation has been reviewed in detail by Griffiths, et al. (1989), and Janesick and Elliot (1991). QEH is not present in the Loral CCDs, because they are front-side illuminated and incorporate MPP operation. This has been verified in component tests at JPL. The absence of QEH means that the devices do not need to be UV-flooded and so decontamination procedures are planned without the constraint of maintaining the UV-flood. 4.7 Flat Field Response The flat field response is uniform within a few percent, with the exception of a manufacturing pattern defect which generates a 3 percent reduction in QE once every 34 rows. This pattern defect is identical in all CCDs. It probably creates an astrometric offset of approximately 3 percent of the pixel height (0.003" in the WFCs) every 34 rows. More precisely, there was a 0.5 micro meter overlap between adjacent 1024xO.5 micro meter raster scans during the construction of the masks used to fabricate the chips. Photometry of point sources imaged onto these defects will be affected, and it will be better to correct flat fields for these rows for such applications. WFPC2 flat fields also include instrumental effects such as vignetting and shadowing by dust particles. Figure 4.4: WFPC2 CCD Flat Field The WFPC2 CCDs have an intrinsically uniform flat field response since they are not thinned, so there are no large-scale chip non-uniformities resulting from the thinning process. MPP operation also improves pixel-to-pixel uniformity because charge transfer is driven deep into the buried n-channel, away from the influence Of Si-SiO2 interface states. The WFPC2 CCD flat fields show an overall pixel-to-pixel response having < 2 percent non-uniformity. Figure 4.4 shows a portion of a WFPC2 CCD flat field obtained during quantum efficiency measurements at JPL. The image illustrates the excellent pixel-to-pixel uniformity of the Loral devices. The 34 row defect is clearly visible, and its amplitude of 3 percent serves to calibrate the gray scale. 4.8 Dark Backgrounds Low dark noise is one of the benefits of MPP, since inverted phase operation suppresses the dominant source of CCD dark noise production (Si-SiO2 surface states). The remaining source of dark noise, thermal generation in the silicon bulk, is determined by the quality of the silicon used in chip fabrication. The intrinsic dark rate of WFPC2 CCDs is < 0.01 e- pixel^-1 s^-1 at temperatures below -80 degrees C. Figure 4.5: Average Dark Rates vs. CCD Row. The temperature set-points for the WFPC2 TEC coolers are: -88, -83, -77, -70, -50, -40, -30 and -20 degrees C. The corresponding approximate median dark rates are given in Table 4.2. For instrument health and safety reasons, GOs cannot command temperature changes. 4.8.1 Sources of Dark Current The dark current appears to have two components: one from electronic sources in the CCD, and a second component whose strength correlates with the cosmic ray flux. The electronic dark current is ~0.001 e- s^-1, consistent with the Thermal Vacuum Test data. The second component of dark current appears only on-orbit, its strength drops towards the edges of each CCD, and it is both chip- and time-dependent. At the current operating temperature, this non-electronic component constitutes up to 80 percent of the total signal measured in the PC. The fraction and overall level are lower in the other chips, and lowest in WF2. This second component ranges from 0.001e- s^-1 (WF2) to 0.005 e- s^-1 (PC). The edge drop of is shown in Figure 4.5, where the average of lines 200-600 for each chip (with hot pixels rejected) is plotted in e- s^-1 as a function of column number. The drop near the edge is consistent with luminescence from the CCD windows, shadowed by a field stop mask just in front of the CCD. Table 4.2: Dark Count Rates. CCD Temperature ( degrees C) Dark count rate (e- s^-1 pixel^-1) ------------------------------------------------------------------------------ -20 10.0 -30 3.0 -40 1.0 -50 0.3 -70 0.03 -77 0.016 -83 0.008 -88 0.0045 ------------------------------------------------------------------------------ A further indication of the possible origin of this second component is the correlation between its amplitude and the cosmic ray activity in the same exposure, as shown in Figure 4.6. For example, the cosmic ray flux in the PC varies from 7x10^5 to 13x10^5 DN per 1000s, while the total dark signal in the PC varies concurrently between 0.0007 and 0.0010 DN s^-1. Similar, though slightly smaller effects are seen in the WFC CCDs. These clues point to cosmic-ray induced scintillation of the MgF2 field-flattening windows as a likely source of the second dark current component. This might be caused by impurities in the MgF2 windows; if so, the window of WF2 must contain substantially less impurities. However, other explanations cannot be completely ruled out at this point. For the great majority of WFPC2 observations, this effect is negligible. In fact, it is noticeable mainly because the true dark rate is very low at the -88 degree C operating temperature. However, observations for which the dark current is an important limiting factor, either due to noise or background flatness, will require special handling to remove the signature of the dark current properly, as its amplitude depends on the time-variable cosmic ray flux. 4.8.2 Darktime As of this writing, the "DARKTIME" keyword in the WFPC2 image headers does not reflect correctly the actual time during which the CCD collects dark current. Instead, DARKTIME is merely set equal to EXPTIME (the exposure time) in the data headers, and this value is used for calibration. The error is small, and usually unimportant, but could be significant for programs aimed at measuring the absolute level of the sky background. The actual darktime in seconds is given by DARKTIME = 60 x int( (t + 13.6)/60 ) + 46.4 + 13.6 x (n - 1) + 60 x (restart) where t is the requested exposure time in seconds, and n is the number of the CCD (PC1=1, WF2=2, etc.), and into indicates the next lower integer. A duration of 13.6s is required both to clear the CCDs before the exposure begins, and to read each CCD after the exposure. External exposures of 180s or longer made with the serial clocks off (CLOCKS=NO; the default setting) suffer an additional 60s of darktime (restart=1). This delay is associated with restarting the serial clocks for readout in exposures where the spacecraft AP-17 processor provides shutter control with loss-of-lock checking. Exposures made with the serial clocks on (CLOCKS=YES) avoid this extra 60s (restart=0). We note that bias frames contain approximately 46.4 + 13.6 x (n - 1) seconds of dark current. As of this writing (April 1996), no attempt is made to subtract this from the bias images when creating calibration files for use in the calibration pipeline. This effect is unimportant for most observations, but could be significant if one averaged many undithered deep exposures of the same field, or if one is interested in measuring the absolute level of the sky background. If the dark current were constant in time, this could be corrected by merely changing the value of DARKTIME used during calibration. However, the hot pixels vary on monthly timescales, so this simple correction is only partially successful. Work is currently underway to improve the way bias calibration files are computed. The timing of dark calibration frames is slightly different from that of external science exposures. Dark calibration frames always have restart=0 in Equation 4.1. Dark calibration frames currently used in the calibration pipeline are averages of ten on-orbit dark frames taken over the space of about 2 weeks. The dark calibration file in the pipeline is revised ~every two weeks to track variations in the hotpixels. As of this writing there are plans to alter this method to both reduce the noise and provide better tracking of hotpixels. The new method would use the average of many (~100) dark frames taken over many months, and then hotpixels would be inserted into the calibration file to correct short-term variations. New files would be generated ~weekly, with only the hotpixels changing from week-to-week. Figure 4.6: Dark Signal vs. Cosmic Ray Flux. Slopes and intercepts ("int") are given on plots. Units are DN/1000s; 1 DN ~ 7 e-. 4.9 Cosmic Rays HST is subjected to cosmic rays and protons from the Earth's radiation belts. The cosmic ray signature in the Loral CCDs is essentially the same as was seen in the WF/PC-1 devices. Electron-hole pairs generated in the thicker substrate by cosmic rays (and infrared photons) are usually removed by recombination in the low resistivity substrate material, because electrons do not diffuse efficiently up to the collecting phase. Cosmic ray events usually deposit significant quantities of charge in more than one pixel. This is due partly to the finite thickness of the CCD detectors, and partly to the less than perfect collection efficiency of each pixel. Figure 4.7 shows a histogram of the number of affected pixels for each cosmic ray event. For the purposes of the figure, a cosmic ray is defined as having a peak pixel value more than 10 DN above the background; and an affected pixel is an attached pixel with a value more than 2 DN above the background. Cosmic ray events do have a clear lower cutoff at around 500 electrons of total signal. The variations in cosmic ray rates caused by orbital position or operating temperature have not been characterized yet. Cosmic ray events impact scientific imaging with WFPC2 in two different ways. Firstly, the WFPC2 CCDs have an epitaxial thickness of about 10 micro meters compared to 8 micro meters for the thinned WF/PC-1 device, and a recombination length of 8-10 micro meters in the substrate. These facts lead to a higher total number of electrons being deposited per event. WFPC2 CCDs also have lower read noise, and so the number of cosmic ray events apparently differs from that of the WF/PC-1 CCDs, since low amplitude events are detected. In practice, this means that the number of pixels apparently contaminated by cosmic rays in an image is higher in WFPC2, although the underlying event rate is similar to that experienced in WF/PC-1. Secondly, stellar images are undersampled and much more difficult to separate from cosmic rays, as is shown in Figure 4.8. Faint stellar images and low energy cosmic rays are indistinguishable. Long science observations are therefore usually broken into at least two exposures (CR-SPLIT) to ensure that events can be identified. Figure 4.7: Histogram of Cosmic Ray Event Sizes. A cosmic ray event is defined by having a peak pixel of at least 10 DN (at gain 7). The average properties of on-orbit cosmic ray events have been determined from examination of several dark exposures, each 2000s long. After bias and dark subtraction, "cosmic rays" were identified in each input frame by first looking for pixels more than 5 sigma above the background, and then including in each event all adjacent pixels more than 2 sigma above the background. Very occasionally, two or more physically separate events will be merged into one as a result of this procedure; visual inspection confirms that in the vast majority of cases, this procedure correctly identifies each event and the area affected by it. The typical value of sigma was 5 to 6 electrons, including both read and dark noise. The region near the borders of each CCD was excluded in order to avoid edge effects, but all results given here are rescaled to the full area of the CCD. One difficulty in this measurement is caused by hot pixels, for some of which the dark current has significant fluctuations from frame to frame; these can be mistakenly identified as cosmic rays when the dark current is at a maximum. Single-pixel events constitute 10 percent of the total number of events identified by our procedure, but at least half of them recur in the same position in several frames, thus identifying them as damaged (hot) pixels, rather than true cosmic rays. Also, unlike the majority of cosmic ray events, single-pixel events tend to have very small total signal; the majority have total signal less than 200 electrons, as expected from hot pixels, while the signal distribution of multiple-pixel events peaks around 1000 electrons. For this reason, single-pixel events have been classified as "bad pixels" rather than "cosmic rays". While we cannot exclude that some true single-pixel events do occur, they are very rare, probably less than 2 percent of the total. Figure 4.8: Comparison of Star Images and Cosmic Ray Events. An 8Ox8O pixel subimage of a 400 second F336W WF2 exposure in omega Cen. Cosmic ray events are frequent, occurring at an average rate of 1.8 events chip^-1 s^-1. The distribution of total signal is shown in Figure 4.9; it has a well-defined maximum at about 1000 electrons, and a cut-off at about 500 electrons. The latter is well above the detection threshold used for the above measurements (25 electrons in the central pixel of the cosmic ray), and is therefore undoubtedly real. Figure 4.9: Histogram of Cosmic Ray Event Energies. The histogram in Figure 4.9 shows the distribution of total energy of all cosmic ray events. One encouraging feature is the very small number of events below about 30 DN. This low energy drop is well above the energy level of excluded single-pixel events. A good approximation to the cumulative distribution of events as a function of total signal is given by a Weibull function with exponent 0.25. This function has the form: N(>S) = N_0exp[-lambda(S^(1/4) - S_0^(1/4))] where N is the total number of events which generate a total signal larger than S. The best fit to the observed events gives N_0=1.4 events chip^-1 s^-1, S_0=700 electrons, and lambda = 0.57 . The fit fails below S_0, and should not be extrapolated to low-signal events. The rate of events with total signal below 700 electrons is 0.4 events chip^-1 s^-1 (i.e. total events per CCD per second is N_0+0.4~1.8). The number of pixels affected by cosmic ray events in a given exposure is a slightly more sensitive function of the threshold used. While there is a clear drop at low signal for both total and peak signal, neighboring pixels can be affected at low levels. Each event (defined as before) affects an average of 6.7 pixels, for about 12 affected pixels chip^-1 s^-1. For a 2000s exposure, this results in about 24,000 affected pixels, or 3.8 percent of all pixels. As cosmic rays are expected to be randomly placed, a pair of such exposures would have about 900 pixels affected in both exposures; cosmic ray correction is impossible for such pixels. For a pair of 1000s exposures, about 220 pixels will be affected in both frames. Cosmic ray activity varies as a function of time, geomagnetic latitude of the spacecraft, and other factors. The average numbers given here are subject to change in individual exposures. After studying about one month's worth of dark exposures, we estimate a total range of about a factor of two in cosmic ray rates. 4.10 Radiation Damage and Hot Pixels In low Earth orbit (LEO) the CCDs are subject to radiation damage from the Earth's radiation belts. The WFPC CCDs are shielded from energetic electrons and about half the incident energetic protons. Long term radiation damage to the CCDs from high energy protons leads to an increase in dark count rate (mainly from the creation of hot pixels), baseline shifts in the CCD amplifiers, and long term degradation of Charge Transfer Efficiency (CTE). There has not been a significant degradation in the amplifier baselines (although occasional small baseline shifts have been seen) and there have been no observed changes in CTE since launch. On the other hand, one of the major radiation damage mechanisms is the creation of new Si-SiO2 interface states, which cause increased dark current rates in affected pixels. In the MPP CCD these states immediately recombine with holes, reducing the gradual increase in dark noise by factors of about 25, compared to normal three-phase CCDs (Woodgate, et al., 1989, Janesick, et al., 1989b). Figure 4.10 is a histogram of the dark current distribution (in e- s^-1) for hot pixels. It contains three curves: solid for the histogram of all hot pixels just before a decontamination (April 7, 1995); dashed only for the pixels that were hot just before the decontamination and were not hot at the beginning of the cycle (March 10); and long-dashed for pixels that were hot at the start of the cycle and were fixed by a decontamination. Thus, the dashed line represents the "new" hot pixels, and the long dashed line represents the fixed hot pixels. The fact that these two curves are very similar shows that the number of hot pixels is roughly in equilibrium. The majority of new hot pixels have low dark current. The hot pixels that constitute the accumulated legacy of previous periods, and thus survived one or more decontaminations, include higher-current pixels. A long term dark frame program is monitoring the increase in the number of hot pixels. The population of hot pixels increases at a rate of approximately 33 pixels CCD^-1 day above a threshold of 0.02 e- pixel^-1 s^-1, while the camera remains at the normal operating temperature. About 80 percent of the new hot pixels return to a normal state at decontamination events when the CCDs are warmed to a temperature of +22 degrees C for 6-12 hours. There is no evidence that the fraction of hot pixels that returns to normal is related to the length of the decontamination. Of those pixels that are not fixed, about half will be fixed after two or three additional decontaminations. After that, the rate of correction decreases. It is conceivable that all hot pixels will be repaired eventually. At the moment there is no evidence of a significant secular increase in the number of hot pixels, and we have a firm upper limit of 8 percent on the fraction of hot pixels that are not repaired after several decontamination cycles. Figure 4.10: Hot Pixel Histogram. In order to deal with the hot pixel problem, we provide monthly lists of possible hot pixels via the World Wide Web and anonymous ftp. Look for hot pixels under WFPC2 Instrument News (/instrument_news/WFPC2/Wfpc2_hotpix via anonymous ftp to stsci.edu). These lists are best used to flag hot pixels as bad. While we do provide an estimate of dark current for each hot pixel as a function of time, there are indications that the noise in hot pixels is much higher than the normal shot noise, and thus dark current subtraction is unlikely to give good results. 4.11 Charge Transfer Efficiency The WFPC2 CCDs have a small but significant charge transfer efficiency (CTE) problem which causes some signal to be lost when charge is transferred down the chip during readout. This has the effect of making objects at higher row numbers (more charge transfers) appear fainter than they would if they were at low row numbers. The effect depends on the temperature of the CCDS. At the original temperature of -76 degrees C, as much as 10-15 percent of the light within a 0.5" radius aperture around a star could be lost for objects at the highest rows. As a result, the CCD operating temperature was changed to -88 degrees C on 23 April 1994. Now the effect seems to have a maximum amplitude of 4 percent for stars with more than 1,500 total detected electrons and no measurable difference in slope is seen for stars up to 20,000 total electrons. However, the brightest stars seem to have a smaller fractional loss. It is also possible that proportional losses for fainter stars, at least with no background, may be larger. Hence, the effect is not well described by either a constant fractional or a constant additive loss per charge transfer. The effect also depends on the amount of background on the chip. There is significantly less CTE effect in the presence of even a moderate (several hundred electrons) background. A major calibration effort is underway to quantify better the dependencies on star color and brightness, and background (see Chapter 8). The photometric calibration of the instrument presented here (and this summary of CTE) is based on Holtzman, et a]. (1995b). It has been crudely corrected for CTE. All of the frames considered are short exposures with essentially no background. For data taken at -88 degrees C, a 4 percent correction ramp was applied to the measured 0.5" radius aperture photometry, in the sense that objects at row 800 were made brighter by 4 percent, but the brightness of objects at the first row was not changed. The correction was applied to bring measurements to the values they would have had in the absence of CTE, or equivalently, the values they would have had if measurements had been made at row 0. Thus, the calibration presented here applies directly to images in which CTE is insignificant (for example, faint stars with high background levels). For images with low background, a CTE correction must be applied before using the calibration presented in this handbook. The current lack of a comprehensive understanding of CTE effects introduces one of the largest uncertainties for WFPC2 photometry. The CTE problems are caused by electron traps in the CCDs which are filled as charge passes through pixels. However, not all traps are accessible to all electrons passing through. Some traps are only accessible if there is significant charge involved. This model suggests that there will not be significant CTE losses in the presence of background, particularly for faint stars, because background electrons fill the traps before such stars pass through. For brighter stars with background there will still be some loss because their charge may access traps that are unaffected by the background that previously clocked through. Faint stars in scenes with little background may suffer from larger losses, although there is no direct evidence for this yet. An apparently related effect surfaced during Cycle 5 as observers compared long and short exposures of the several stellar fields -- there was evidence for a discrepancy between the photometric zeropoint of long and short exposures. Further investigation shows that "long vs. short" is probably a misnomer. The level of background appears to be the main parameter rather than the exposure time. The magnitude difference measured between short and long exposures is more pronounced for faint stars in large apertures, where it can reach 0.05 mag, and is essentially absent for stars with more than 1000 total counts. The dependence on aperture and magnitude appears consistent with a charge transfer efficiency problem. The offset of faint star magnitudes can be explained by a loss of 0.3 DN (2 e-) in each pixel used in the aperture. Further testing and analysis of this effect are underway. A side effect of the CTE problem, is that residual trails can appear in exposures after a highly exposed image is read out (Figure 4.11). The trail is caused by charge which is trapped during read out of the highly exposed image, which is then slowly released during subsequent exposures. The effect is most pronounced when long exposures in through low throughput filters (narrow band or UV filters) immediately follow a highly exposed image (usually a broad band filter). 4.12 Read Noise and Gain Settings The CCDs and their associated signal chains have readout noise levels (in the absence of signal shot noise or interference) of approximately 5e-. The analog-to-digital converter is highly accurate, and makes virtually no contribution to the read noise, other than the normal information loss caused by digitization of the signal. The conversion factors from detected electrons (QE x number of incident photons) to data numbers (DN) are tabulated in Table 4.3, as are read noise and linearity ("gamma" is the power law index relating detected DN to input flux). Note that all calculations of sensitivity in this manual assume gains of 7 and 14 for two gain channels, choices very close to the measured gains. The photometric calibration is based on an assumed exact gain of 14 in all CCDS. The measurements given here were derived from thermal vacuum testing. On-orbit measurements have confirmed that the gain ratios are correct to within a possible systematic error of 1 percent-which will feed directly into a photometric calibration error for gain 7 data, as most of the photometric calibration was done with gain 14 data. Note that the gain ratios are known much more accurately than the individual gains; they are derived from flat field ratios instead. Also, note that the Phase II proposal instructions refer the ~14 e- DN^-1 setting as ATD-GAIN=15. Table 4.3: Signal Chain Gains. Parameter Gain PC1 WF2 WF3 WF4 ------------------------------------------------------------------------------ Noise "7" 5.24 +- 0.30 5.51 +- 0.37 5.22 +- 0.28 5.19 +- 0.36 "15" 7.02 +- 0.41 7.84 +- 0.46 6.99 +- 0.38 8.32 +- 0.46 Gain "7" 7.12 +- 0.41 7.12 +- 0.41 6.90 +- 0.32 7.10 +- 0.39 "15" 13.99 +- 0.63 14.50 +- 0.77 13.95 +- 0.63 13.95 +- 0.70 Gamma "7" 1.0015+-0.0006 1.0015+-0.0006 1.0020+-0.0006 1.0038+-0.0007 "15" 1.0004+-0.0001 1.0023+-0.0004 1.0032+-0.0006 1.0018+-0.0012 14/7 ratio 1.987 +- 0.02 2.003 +- 0.02 2.006 +- 0.02 1.955 +- 0.02 ------------------------------------------------------------------------------ Figure 4.11: Images Illustrating CTE Residual Trail. (a) Image of star field taken in broad band filter on PC1. (b) 1500s dark exposure taken immediately following (a). Read out direction is towards bottom of image. Cosmic rays have been removed. ------------------------------------------------------------------------------ CHAPTER 5: Point Spread Function In This Chapter... Effects of OTA Spherical Aberration Aberration Correction Wavefront Quality CCD Pixel Response Function Model PSFs PSF Variations with Field Position PSF Variations with Time Large Angle Scattering Ghost Images Optical Distortion 5.1 Effects of OTA Spherical Aberration The OTA spherical aberration produces a Point Spread Function (PSF-the apparent surface brightness profile of a point source), as presented to the instruments, with broad wings. Briefly, the fraction of the light within the central 0.1" was reduced by a factor of about 5. The resulting PSF had "wings" which extended to large radii (several arcseconds), greatly reducing the contrast of the images and degrading the measurements of sources near bright objects or in crowded fields. Burrows et al. (1991, Ap. J. Lett. 369, L21) provide a more complete description of the aberrated HST PSF. Figure 5.1 shows the PSF in three cases. It shows the aberrated HST PSF, the WFPC2 PSF, and for comparison the PSF that would be obtained from a long integration if HST were installed at a ground based observatory with one arcsecond seeing. All of the PSFs were computed at 4000A. The FWHM of the image both before and after the installation of WFPC2 is approximately proportional to wavelength, at least before detector resolution and MTF effects are considered. (The WF/PC-1 core was approximately 50 percent broader than the core that is obtained with WFPC2). Figure 5.2 shows the encircled energy (EE), the proportion of the total energy from a point source within a given radius of the image center, for the same three cases. The WFPC2 curve shown is the average of measurements taken with F336W and F439W. It can be seen that the core of the image in WFPC2 contains most of the light. At this wavelength, 65 percent of the light is contained within a circle of radius 0.1". However, this proportion is considerably less than the optics deliver. The reason for this is discussed in the Section 5.4. Encircled energy curves for other filters are shown in Figures 5.3 and 5.4; note that these curves are normalized to unity at 1.0" radius. Figure 5.1: PSF Surface Brightness. The percentage of the total flux at 4000A failing on a PC pixel as a function of the distance from the peak of a star image. 5.2 Aberration Correction WFPC2 has corrective figures on the relay secondary mirrors where the primary mirror is imaged; this optical correction recovers near-diffraction limited images over the entire CCD fields-of-view. Proper correction requires tight optical alignment tolerances, which are facilitated on-orbit by actuated optics. The corrective optics enable essentially all of the scientific objectives of the original WF/PC-1 to be met. Figure 5.2: Encircled Energy. The percentage of the total flux at 4000A within a given radius of the image peak. Table 5.1: Wavefront Error Budget. Camera WFC(F/12.9) PC(F/28.3) ------------------------------------------------------------------------------ Design error lambda/143 lambda/50 Fabrication and alignment error lambda/14.7 lambda/14.7 Alignment stability error lambda/25 lambda/25 Total wavefront error lambda/12.6 lambda/12.3 ------------------------------------------------------------------------------ Through a number of independent analyses, based on investigations of star images obtained on-orbit, and the examination of fixtures used during the figuring of the primary mirror, the aberrations of the HST optics were accurately characterized. The primary mirror was figured to an incorrect conic constant: -1.0139-+0.005 rather than the -1.0023 design requirement, resulting in a large amount of spherical aberration. The optical design of WFPC2 creates an image of the OTA primary mirror near the surface of the relay Cassegrain secondary mirror in each of its channels. This design minimizes vignetting in the relay optics, but more importantly, facilitates correction of spherical aberration in the OTA primary by application of the same error (but with opposite sign) to the relay secondary. The optical figure of the WFPC2 secondary mirrors contains a compensating "error" in the conic constant. By adopting a prescription within the error bars for the HST primary mirror, corrective secondary mirrors were made with sufficient accuracy that the residual spherical aberration in the WFPC2 wavefront is small compared to other terms in the WFPC2 optical wavefront budget. Figure 5.3: Encircled Energy for CCD PC1. The fraction of energy encircled is plotted vs. aperture radius for several filters. Curves are normalized to unity at a radius of 1.0". From Holtzman, et al., 1995a. Figure 5.4: Encircled Energy for CCD WF3. The fraction of energy encircled is plotted vs. aperture radius for several filters. Curves are normalized to unity at a radius of 1.0". From Holtzman, et al., 1995a. On the other hand, new and stringent alignment requirements were created by the steep optical figure on the corrective relay secondary mirrors. The primary mirror image must be accurately centered on the corrective mirror, and must have the correct magnification. Centering is the most demanding requirement. A failure to center accurately would create a new aberration in the form of coma. A misalignment of 7 percent of the pupil diameter introduces as much RMS wavefront error as was present in the form of spherical aberration prior to the introduction of corrective optics. The new requirements for alignment accuracy and stability led to the introduction of a tip-tilt mechanism on the pick-off mirror, to compensate for camera alignment uncertainties with respect to the OTA, and actuated fold mirrors which can compensate for internal misalignments. There was an additional term in the CEIS specification of the overall instrument wavefront error budget for alignment stability. It is lambda/25 RMS at 6328A, as shown in Table 5. 1. "Design error" refers to the aberrations inherent in the design itself, which would be seen if the optics conformed perfectly to their specifications. All of the optics were fabricated and integrated into the WFPC2 optical bench. It was established on the basis of component tests, end-to-end optical interferometry, and through focus phase retrieval, that the WFPC2 optical system performed within the stated tolerances for "fabrication and alignment" in the laboratory environment. What remained was to demonstrate the stability of the optical alignment after launch vibration and in response to the thermal environment on-orbit. The "stability" line anticipated these uncertainties, and has been verified during early science operations. 5.3 Wavefront Quality The conclusion of the extensive optical testing in thermal vacuum was that the cameras are well corrected to within the specifications. The measured wavefront errors in the four cameras were 1/30, 1/17, 1/40, and 1/21 waves at 6328A. The dominant problem was a small difference in focus between the four cameras. The actuated fold mirrors and pick-off mirror mechanism performed flawlessly in correcting residual coma aberrations in the image, and enabled the on-orbit alignment procedures. By defocusing the OTA secondary mirror, a very accurate alignment of the cameras was accomplished. A side effect was that the aberrations in each camera were measured (Krist and Burrows, Applied Optics, 1995). The results are given in Table 5.2. These values were used in generating the simulated PSFs given in section 5.5, "Model PSFS", on page 92. Table 5.2: Aberrations in Each Camera. The numbers quoted are RMS wavefront errors in microns over the HST aperture (Zernike coefficients). Aberration PC1 WF2 WF3 WF4 ------------------------------------------------------------------------------ Z4 Defocus 0.0000 0.0410 0.0640 0.0480 Z5 0 degrees Astig 0.0229 0.0109 0.0126 0.0163 Z6 45 degrees Astig 0.0105 0.0041 0.0113 0.0190 Z7 V2 Coma 0.0000 0.0012 -0.0037 -0.0090 Z8 V3 Coma -0.0082 0.0061 -0.0100 0.0019 Z9 X Clover 0.0063 0.0121 0.0010 0.0096 zio Y Clover 0.0023 0.0091 0.0130 0.0042 Zil Spherical -0.0131 -0.0215 -0.0265 -0.0247 Z22 5th Spherical 0.0034 0.0034 0.0036 0.0029 Zonal Errors 0.0180 0.0180 0.0180 0.0180 ------------------------------------------------------------------------------ Total (RS5) 0.0353 0.0537 0.0755 0.0637 ------------------------------------------------------------------------------ Budget 0.0813 0.0794 0.0794 0.0794 ------------------------------------------------------------------------------ 5.4 CCD Pixel Response Function From Thermal Vacuum testing, there was evidence that the images are not as sharp as expected, despite the good wavefront quality. The decrease in sharpness corresponds to a loss in limiting magnitude of about 0.5 magnitudes in the WF cameras, and less in the PC. Further testing, by covering a flight spare CCD with a 2 micro meter pinhole grid in an opaque metallic mask and illuminating it with a flat field source, showed that even when a pinhole was centered over a pixel only about 70 percent of the light was detected in that pixel. For practical purposes, the effect can be modeled as equivalent to about 40 mas RMS gaussian jitter in the WFC, and 18 mas in the PC (as compared with the typical real pointing jitter of ~3 mas delivered by the excellent HST pointing control system). Alternatively, at least in the V band, it can be modeled by convolving a simulated image by the following kernel, which gives the pixel response function averaged within pixels: 0.016 0.067 0.016 K = 0.080 0.635 0.080 0.015 0.078 0.015 One clue is the wavelength dependence of the observed sharpness: the results from the 2 micro meter pinhole grid test get worse at longer wavelengths. This may reflect the greater penetration into the silicon of low energy photons, which facilitates the diffusion of photoelectrons across the pixel boundaries defined by the frontside gate structure. There is also evidence for sub-pixel QE variations at the 10 percent level. There is an implied dependence on pixel phase for stellar photometry. This has been seen at about the 1 percent level in on-orbit data. The work of Jorden, Deltom, and Oates (Greenwich Observatory Newsletter 9/93) has yielded quite similar results, and suggests that sub-pixel response must be taken into account when seeking to understand the behavior of all CCD detectors forming undersampled images. 5.5 Model PSFs Considerable effort has gone into the modeling of the HST point spread function (PSF), both in order to measure the optical aberrations in support of the WFPC2, COSTAR, and advanced scientific instruments, and to provide PSFs for image deconvolution in the aberrated telescope. Such PSFs are noise free and do not require valuable HST observing time. Software to generate model PSFs for any filter and at any location within the field-of-view is available from the STScI TIM package, Hasan and Burrows 1993; TinyTIM package, Krist 1995). The results are illustrated in Tables 5.3 and 5.4 for the PC1 and WF2 cameras, respectively. A representative PSF is on the left in each panel. It meets the wavefront error budget, with the measured mix of focus, coma, astigmatism, and spherical aberration. It has been degraded by the pixel response function as discussed in the Section 5.4. On the right is the diffraction limited case for comparison. In each case the percentage of the total flux in a central 5x5 pixel region of a point source is displayed. The peak of the star image can be at an arbitrary point relative to the boundaries of the CCD pixels. Two cases are shown: one where the star is approximately centered on a pixel, and one where it is approximately centered at a pixel comer. As a consequence of the under-sampling in the WFPC2, the limiting magnitude attainable in the background limit varies by about 0.5 magnitude, depending on the position of the source within the CCD pixel. This point is discussed in more detail in Chapter 6. Neither observed nor modeled PSFs will provide a perfect match to the PSF in actual science observations, due to modeling uncertainties, the "jitter" in the HST pointing, and orbit to orbit variations in telescope focus ("breathing"-which seems to be generally limited to about 1/20 wave peak-to-peak). Jitter is not predictable but can be recovered to a reasonable extent for observations obtained in Fine Lock. In long exposures, up to about 10 mas of apparent pointing drift may occur as a result of the breathing effects in the FGS, although smaller variations of ~3 mas are typical. Table 5.3: PC Point Spread Functions. Shown as percentages (out of 100 percent) of the total flux in a 5 by 5 pixel region. On the left in each case is a model PSF with the observed wavefront errors and pixel response function. On the right is the diffraction limited case for comparison. WFPC2 Model PSF 2000 A: Peak near Diffraction Limited PSF corner of PC pixel 0.9 2.3 2.3 0.9 0.3 0.2 0.6 0.5 0.3 0.1 2.5 10.3 12.7 2.5 0.5 0.6 17.6 20.9 0.5 0.1 1.9 11.2 13.3 2.6 0.5 0.5 20.9 26.0 0.6 0.2 0.9 2.1 2.7 1.4 0.3 0.3 0.5 0.6 0.3 0.1 0.3 0.5 0.5 0,4 0.2 0.1 0.1 0.2 0.1 0.2 Peak near center of pixel 0.3 0.7 1.3 0.8 0.5 0.1 0.3 0.4 0.2 0.2 0.7 2.6 6.4 3.2 0.8 0.3 0.6 4.9 0.6 0.3 1.2 6.3 25.0 6.9 1.4 0.4 4.9 62.9 6.3 0.4 0.6 2.2 5.9 3.6 0.9 0.3 0.5 6.3 0.7 0.4 0.4 0.8 1.4 1.1 0.4 0.2 0.3 0.4 0.4 0.2 4000 A: Peak near corner of PC pixel 0.9 2.5 3.5 1.2 0.2 0.3 2.5 2.5 0.3 0.1 3.6 10.8 12.0 3.1 0.4 2.5 14.3 15.8 2.7 0.2 2.7 11.5 12.9 3.5 0.4 2.5 15.8 17.6 3.0 0.2 0.8 3.0 3.4 1.3 0.3 0.3 2.7 3.0 0.4 0.1 0.2 0.4 0.4 0.3 0.2 0.1 0.2 0.2 0.1 0.1 Peak near center of pixel 0.3 0.6 0.8 0.9 0.3 0.1 0.2 0.3 0.2 0.1 0.9 4.0 6.6 4.9 0.8 0.2 3.8 4.4 3.8 0.2 0.9 6.9 26.1 7.0 0.8 0.3 4.5 49.4 5.1 0.4 0.5 3.4 6.8 4.9 0.9 0.2 3.8 5.0 4.3 0.2 0.2 0.6 0.8 0.9 0.4 0.1 0.2 0.4 0.2 0.1 6000 A: Peak near corner of PC pixel 2.0 2.6 3.4 2.4 0.5 2.1 2.3 2.0 2.1 0.2 3.4 9.8 10.4 2.9 0.6 2.2 11.4 12.8 2.1 0.6 2.8 10.6 11.2 3.2 0.6 2.0 12.9 14.1 2.1 0.6 1.6 2.8 3.1 2.4 0.5 2.0 2.0 2.1 2.3 0.2 0.4 0.6 0.7 0.5 0.2 0.2 0.5 0.6 0.2 0.1 Peak near center of pixel 0.5 1.2 1.7 1.7 0.6 0.2 1.5 1.8 1.4 0.3 1.7 3.1 5.9 3.6 1.6 1.5 2.4 4.3 2.2 1.6 2.0 6.0 20.7 6.6 1.8 1.8 4.4 31.6 5.3 2.0 1.2 2.9 6.2 3.7 1.7 1.4 2.1 5.4 2.3 1.7 0.4 1.3 2.0 1.7 0.6 0.2 1.5 1.9 1.7 0.3 8000 A: Peak near corner of PC pixel 1.6 1.9 2.2 1.9 1.1 1.8 0.9 0.9 1.5 1.1 2.1 9.3 9.7 2.1 1.1 0.9 11.7 12.6 1.0 1.4 2.0 9.8 10.1 2.4 1.1 0.9 12.6 13.3 1.0 1.5 1.4 2.1 2.1 1.8 1.1 1.5 1.0 1.0 1.6 1.2 0.8 1.2 1.3 1.1 0.4 1.1 1.4 1.5 1.2 0.2 Peak near center of pixel 1.2 1.4 1.5 1.8 1.4 1.3 1.8 1.1 1.6 1.4 1.8 2.5 6.0 2.8 1.6 1.8 1.5 6.2 1.7 1.6 1.6 6.0 15.4 6.6 1.5 1.1 6.2 22.5 7.1 1.0 1.3 2.7 6.3 3.0 1.7 1.6 1.7 7.1 1.9 1.7 1.0 1.4 1.5 1.7 1.4 1.4 1.6 1.0 1.7 1.5 WFPC2 Model PSF 2000 A: Peak near Diffraction Limited PSF corner of PC pixel 0.5 1.8 2.1 0.8 0.3 0.4 0.4 0.4 0.3 0.1 1.5 9.9 13.6 3.1 0.5 0.4 15.4 21.5 0.4 0.1 1.5 9.8 24.0 4.7 0.5 0.4 21.5 33.4 0.4 0.1 0.5 2.1 3.6 1.3 0.3 0.3 0.4 0.4 0.4 0.1 0.2 0.4 0.4 0.2 0.2 0.1 0.1 0.1 0.1 0.1 Peak near center of pixel 0.2 0.4 0.5 0.5 0.3 0.2 0.2 0.2 0.1 0.2 0.3 1.7 5.3 3.0 0.8 0.1 0.8 1.3 0.6 0.1 0.5 5.4 28.5 14.0 2.2 0.2 1.3 86.2 1.4 0.2 0.3 2.2 9.1 4.5 1.0 0.1 0.6 1.4 0.9 0.2 0.3 0.6 1.2 0.8 0.3 0.2 0.1 0.2 0.2 0.2 400o A: Peak near corner of WF pixel 0.8 2.6 2.5 0.8 0.2 0.4 0.5 0.5 0.2 0.1 2.6 16.1 16.4 3.0 0.4 0.5 17.7 21.2 0.6 0.1 2.0 12.9 17.3 3.0 0.4 0.5 21.2 26.7 0.6 0.1 0.6 2.0 2.6 0.9 0.2 0.2 0.6 0.6 0.4 0.2 0.2 0.4 0.3 0.2 0.1 0.1 0.1 0.1 0.2 0.2 Peak near center of pixel 0.3 0.6 0.7 0.5 0.3 0.3 0.2 0.3 0.2 0.2 0.6 3.1 7.2 3.4 0.7 0.2 0.7 3.7 0.7 0.2 0.9 8.5 33.3 10.2 1.2 0.3 3.7 68.8 5.3 0.4 0.4 2.5 8.8 3.0 0.6 0.2 0.7 5.3 0@8 0.2 0.2 0.5 0.9 0.5 0.3 0.2 0.2 0.4 0.2 0.3 6000 A: Peak near corner of WF pixel 0.7 2.6 2.6 0.9 0.3 0.2 0.5 0.4 0.2 0.2 3.0 14.9 15.6 3.3 0.5 0.5 18.3 20.7 0.5 0.3 2.2 13.7 16.3 3.2 0.4 0.4 20.7 24.2 0.6 0.2 0.6 2.3 2.9 0.7 0.2 0.2 0.5 0.6 0.2 0.2 0.2 0.3 0.3 0.2 0.2 0.2 0.3 0.2 0.2 0.2 Peak near center of pixel 0.3 0.7 0.9 0.6 0.2 0.2 0.3 0.2 0.3 0.2 0.7 4.1 7.8 4.1 0.8 0.3 1.8 6.1 1.9 0.3 1.0 8.6 30.4 9.4 1.3 0.2 6.1 54.9 6.2 0.3 0.5 2.8 8.2 3.5 0.6 0.3 1.9 6.2 2.5 0.3 0.2 0.5 1.0 0.6 0.3 0.2 0.3 0.3 0.3 0.2 8000 A; Peak near corner of WF pixel 1.0 3.0 2.9 1.1 0.3 0.1 2.0 2.0 0.2 0.2 3.6 13.1 13.6 4.0 0.5 2.0 15.8 17.3 2.2 0.1 2.6 12.6 14.2 3.5 0.5 2.0 17.3 19.3 2.5 0.1 0.8 3.2 3.6 1.0 0.3 0.2 2.2 2.5 0.2 0.2 0.2 0.4 0.5 0.2 0.1 0.2 0.1 0.1 0.2 0.1 Peak near center of pixel 0.2 0.8 0.8 0.7 0.3 0.1 0.2 0.3 0.2 0.1 0.9 4.6 6.9 4.4 1.0 0.2 3.5 4.6 3.4 0.2 0.9 7.5 30.8 8.1 1.2 0.3 4.6 52.0 5.0 0.3 0.5 3.1 7.3 3.7 0.7 0.1 3.4 5.1 3.9 0.2 0.2 0.5 1.0 0.7 0.2 0.1 0.2 0.3 0.2 0.1 5.6 PSF Variations with Field Position The WFPC2 PSFs vary with field position due to field-dependent aberrations, obscuration shifting, and scattering. This will complicate photometry, PSF subtraction, and deconvolution (Krist, 1995). The coma and astigmatism aberrations vary significantly within a camera across the field-of-view. These variations are simply part of the optical design. At the extreme comers of the WFC CCDS, away from the OTA axis, there is about 1/5 wave of astigmatism (referenced at 633 nm), which decreases to nearly zero at the CCD centers. Astigmatism at this level causes the PSF core to become elliptical and slightly less sharp; note flattening of PSF at pixels positions (54,777) and (605,148) in Figure 5.5. Coma also varies, but to a much lesser extent. Coma and astigmatism variations are considerably smaller in PC1. Figure 5.5: PSF Variations with Field Position - Aberrations. Nine observed PSFs (filter F814W) are shown from a widely spaced grid on WF3. CCD pixel position are labeled. Note flattening of PSF in the (54,777) and (605,148) positions. The obscuration patterns due to the camera optics (relay secondary mirror and spiders) appear to shift with respect to the OTA obscurations, depending on field position. The interacting diffraction patterns of the WFPC2 and OTA spiders cause ripples in the spider diffraction spikes, which vary with field position as the two spiders shift relative to each other. In Figure 5.6 the OTA spider is hidden behind the WFPC2 spider at the field center and hence the diffraction spikes there have a simple, smooth appearance (c.f. position 446,425). At the CCD comers, however, one or more vanes of the OTA spider moves out from behind the WFPC2 spider, and the double set of obscurations causes a "beating" pattern in the diffraction spikes. The spiders also interact with light diffracted from zonal errors in the OTA mirrors, causing streaks in the scattering halo which vary in position and intensity. Figure 5.6: PSF Variations with Field Position - Obscuration Shifts. Five saturated PSFs observed in F814W are shown from a widely spaced grid on WF4. Note changes in spider diff raction spikes. CCD pixel position are labeled. 5.7 PSF Variations with Time The shape and width of observed PSFs varies slightly over time, due to the change in focus of the telescope. The focus variation consists of two terms: a secular change, due to the ongoing shrinkage of the Metering Truss Assembly, at a currently estimated rate of 0.85 micro meters month^-1, and short-term variations, typically on an orbital time-scale (the so-called "breathing" of the telescope). The breathing is probably due to changes in the thermal environment as the telescope moves through its orbit, and has a typical peak-to-peak amplitude of 4 micro meters; larger variations are occasionally seen. These small focus shifts will impact photometry performed with small (few pixel radius) apertures. Typical +- 2 micro meters focus shifts will result in photometric variations in the PC1 of 6.8 percent, 4.5 percent, 2.0 percent, and 0.2 percent for aperture radii of 1, 2, 3 and 5 pixels, respectively in F555W. This is based on the focus monitoring data taken over the period from July 1994 to January 1996 (Figure 5.7). Hence, "breathing" is often one of the major sources of errors for small-aperture photometry. However, relative photometry (i.e. the difference in magnitudes of stars in the same image) is not affected by this variation very much since all the stars in an image are impacted by the defocusing in a similar way. The variations are, of course, smaller for the WF chips because the pixels are roughly twice as large. Figure 5.7: Measured OTA Focus Position (microns) as Function of Number of Days since January 1, 1994. Focus position is defined as difference between the optimal position of OTA secondary mirror and actual position, in microns. Figure 5.8: Measured Aperture Correction, V(r) - V(r=10 pix), in Magnitudes as Function of Shift from Optimal Focus. Data are given for aperture radii r=1, 2, 3, and 5 pixels for F555W filter on CCD PC1. Systematic errors due to the secular focus drift can be corrected using the aperture corrections phased with the focus change (see Figure 5.8): ap_corr = ap_corr_nominal + a(r) x d where ap_corr_nominal is the nominal aperture correction (mag) as derived from Table 2a in Holtzman et al. (1995a), a(r) is the flux variation per 1 micro meter of focus drift (mag per micron) using an aperture with radius r (pixels), and d (micro meter) is the focus shift from the nominal position. The monitoring data mentioned above yield for PC1, F555W, the following values for a(r): a(1 pix) = 0.0338 +- 0.0038 a(2 pix) = 0.0226 +- 0.0024 a(3 pix) = 0.0100 +- 0.0018 a(5 pix) = 0.00105 +- 0.0015 Values for other filters and other chips will be compiled in an instrument science report later in 1996. Large focus changes, with amplitudes up to 10 micro meters, are seen occasionally (See Hasan and Bely, Restoration of HST Images and Spectra II, p. 157). On May 1, 1994, and February 27, 1995, a short-lived defocusing of the telescope of up to logm was seen, probably due to extreme thermal conditions. Such a defocusing causes an increase of the PSF width by about 5-10 percent and a significant change in its shape, especially evident in the PC both because of its higher resolution and its astigmatism (which makes the out-of-focus image appear elongated). The change in the PSF appears to be modeled adequately by the TIM software. (See Hasan and Bely, Restoration of HST Images and Spectra II, p. 157). (See sample PSF subtraction in Figure 7.2 on page 148). 5.8 Large Angle Scattering Analysis of the WFPC2 saturated star images indicate that the large angle scattering (> 3" from a star) is significantly higher than expected. Three data sets were used to determine the WFPC2 scattering. The first set was from the SMOV Ghost Check proposal 5615, in which 100 second images of delta Cas (V=2.7) were obtained at the center of each chip in F502N. The second set was a series of 6 second exposures of Vega (V=O.O) centered on WF2 through F410M (WFPC2 GTO proposal 5205). The third set was epsilon Eridani (V=3.73) centered on the PC and taken through F631N (500 sec each) and F953N (2200 sec each). These were from GTO proposal 5611. WFPC2 scattering was determined by computing the azimuthal average and azimuthal median profiles. The regions near the diffraction spikes and saturated columns were not used. The profiles were determined using images corrected for horizontal smearing. After renormalizing to the XCAL fluxes the profiles agreed. The measurements indicate that the average scatter in WFPC2 is an order of magnitude greater than in WF/PC-1. The increase is due to scattering in WFPC2, not the OTA. In the WFPC2 images, the pyramid edge shadow is not visible in the scattered light; the light is spread out to the chip edges, indicating that most of the scattering occurs after the pyramid. However, the light level in adjacent channels is back down at the WF/PC-1 levels as shown in Figure 5.9. The scattering does not show any strong dependence on wavelength between 410 nm and 953 nm, within the uncertainties of the measurements. The scattered light is not uniform. There are high frequency spatial structures in the form of streaks radiating outwards from the star. These features are probably both wavelength and position dependent, and so cannot be readily subtracted. The source of the WFPC2 scattering may be the CCDS. The WF/PC-1 CCDs were back illuminated and had shiny surfaces. The electrode structure was not visible over most of the wavelength range. The WFPC2 CCDS, however, are front illuminated, so the electrode structure is visible and may be scattering the light There was a large ghost in WF/PC-1 due to a reflection between the CCD and filter, but no such feature has been seen in WFPC2. The flux from this missing ghost may instead constitute part of the scatter. (See also related material in Observing Faint Targets Near Bright Objects" on page 146.) Figure 5.9: Large Angle Scattering. The proportion of the total flux in F555W falling per square arcsecond as a function of the distance from the peak of a saturated stellar image. These curves are for a target in the PC. Note the large drop in the scattered light level when looking in an adjacent camera. 5.9 Ghost Images Common ghost images result from internal reflections in the filters and in the field-flatteners. Two filter ghosts, caused by double (and quadruple) reflection inside the filter, are visible below and to the right of the star in Figure 5.10. The position and brightness of these ghosts varies from filter to filter, typically being most obvious in interference filters. The comatic shape of the ghost is caused by the camera optics being effectively misaligned for the light path followed by the ghost. The relative position of these ghosts does not vary much over the field. An additional ghost is caused by an internal reflection inside the MgF2 field flattener lens immediately in front of each CCD (Figure 5.11). The field flattener ghost is doughnut shaped (image of OTA pupil) in the WFC, but is smaller and more disk-like on the PC. This ghost contains ~0.15 percent of the total energy of the star. It is positioned on a line through the CCD center and the bright star; the distance from the ghost to the CCD center is 1.25 to 1.4 times the distance from the bright star to the CCD center. This geometry results from curvature of the field flattener lens. The large ghost image expected to be caused by reflection off the CCD back to the filter and then back to the CCD is not seen. It was deliberately eliminated in the PC by tilting the CCD slightly. Figure 5.10: Saturated Stellar Image Showing Filter Ghosts. Intensity scale is logarithmic. Figure 5.11: Saturated Stellar Image Showing Field Flattener Ghost on WF2. 5.10 Optical Distortion The WFPC2 cameras have significant geometric distortion which not only affects astrometry, but also affects photometry (because the extended sources used to generate flat fields have an induced change in apparent surface brightness). It can be a large effect, with "true" positions differing from observed positions by several pixels in the comers of the cameras. The distortion is wavelength dependent in the ultraviolet, because it is partially caused by the MgF2 field flattener in front of each CCD. It is sensibly wavelength independent in the visible. An estimate of the geometric distortion in the WFPC2 cameras was made (Holtzman et al., PASP 107, 156) from a series of 30 F555W exposures in a dense stellar field. Between each of these exposures, the telescope pointing was shifted by 16" in a 6 x 5 grid along rows and columns. On any given frame, positions of several dozens of stars were measured. Any pair of stars which appeared on three or more of the 30 frames was used to do a cubic distortion solution. The cubic distortion coefficients are of the form: x_corr = C_1 + C_2x + C_3y + C_4x^2 + C_5xy + C_6y^2 + C_7x^3+ C_8x^2y + C_9xy^2 + C_10^3 y_corr = D_1 + D_2x + D_3y + D_4x^2 + D_5xy + D_6y^2 + D_7x^3+ D_8x^2y + D_9xy^2 + D_10^3 The coefficients are given in Table 5.5. The input (xy) in the above equation are offsets from the center of each CCD in pixel units: x = x_obs - 400 y = y_obs - 400 where (x_ob, y_obs) are the "observed" pixel positions on each CCD. The corrected values (x_corr y_corr) are in a system with the origin near the pyramid apex, and the units are PC1 pixels. Hence PC1 in is quadrant 1, WF2 in quadrant 2, etc. Application of the transformation brings positions of all chips into the orientation of Pc1. The pixel scale can be estimated from the commanded offsets between the frames (relying on the FGS scale and distortion calibrations). It comes out as 0.04554 +- 0.00001 arcseconds pixel^-1 in the PC, and hence 0.09961, 0.09958, and 0.09964 arcseconds pixel^-2 in WF2, 3 and 4 respectively. An independent check on an astrometric standard field (M67) yielded 0.04555 arcseconds pixel^-1 in the PC. These plate scales refer to the scale at the center of the chip in filter F555W. The true scale is lower elsewhere on the chip because of distortion, and there is some wavelength dependence in the scale even for visible wavelengths. Table 5.5: Cubic Distortion Coefficients. Coefficient PC1 WF2 WF3 WF4 ------------------------------------------------------------------------------ C1 3.54356E+02 -8.12003E+02 -8.07068E+02 7.72904E+02 C2 1.00021 E+00 1.95172E-02 -2.18757E+00 1.37619E-02 C3 9.79758E-04 -2.18805E+00 -8.51284E-03 2.18899E+00 C4 9.84222E-08 -1.51015E-06 -3.21458E-07 1.33326E-06 C5 -6.31327E-09 5.06693E-06 2.77459E-06 -4.00828E-06 C6 -7.19983E-07 1.04063E-06 -2.06874E-06 -5.71282E-07 C7 -3.73922E-08 -4.55094E-10 7.39223E-08 -2.00520E-09 C8 6.651OIE-10 7.50249E-08 -1.18833E-09 -7.86673E-08 C9 -3.51470E-08 -1.51652E-09 7.73056E-08 -3.37491E-09 C10 -2.55003E-09 7.34568E-08 -3.72971E-10 -7.81490E-08 D1 3.43646E+02 7.66592E+02 -7.71489E+02 -7.74638E+02 D2 1.00544E-03 2.18607E+00 7.40319E-03 -2.18688E+00 D3 9.99786E-01 2.02006E-02 -2.18554E+00 1.29831E-02 D4 -5.80870E-07 -3.79469E-06 -1.49841E-06 1.30814E-06 D5 -5.07221E-07 -2.99518E-06 2.96650E-06 2.01760E-06 D6 3.41574E-07 1.05789E-07 -2.02756E-07 -8.47656E-07 D7 -1.08091E-09 -7.47857E-08 -1.45720E-09 7.70408E-08 D8 -3.42070E-08 -1.67666E-09 7.72421E-08 -1.68514E-09 D9 1. 17819E-09 -7.55302E-08 8.52633E-10 7.75257E-08 ------------------------------------------------------------------------------ The cubic distortion coefficients can be used to derive effective pixel areas as presented in Figure 5.12. Contours are shown at half percent levels. Measurements of total brightness or total counts (as opposed to measurements of surface brightness) should be corrected by multiplying the science image by Figure 5.12. This correction image is also available in the HST data archive as file f1k1552bu.r9h. Figure 5.12: Integrated Photometry Correction Induced by Camera Distortions. ------------------------------------------------------------------------------ CHAPTER 6: System Throughput and SNR / Exposure Time Estimation In This Chapter... System Throughput On-Line Exposure Time Calculator Target Count Rates Sky Background Signal-to-Noise Ratio Estimation Exposure Time Estimation Sample SNR Calculations Red Leaks in UV Filters Time Dependence of UV Response 6.1 System Throughput A decision on a suitable exposure time will require the combination of * The overall spectral response of the system (Figure 2.4). * The spectral transmission of the filters (Chapter 3 and Appendix A). * The spectral energy distribution and spatial profile of the target. * The point response function and pixel size of the instrument (Chapter 5). * Criteria for specifying desirable charge levels. When the transmissions of filters T(lambda) are combined with the overall system response Q(lambda), we obtain detective quantum efficiency plots (electrons-per-photon as a function of lambda) for each filter. These DQE plots link the output of the CCD to the photon flux at the input to an unobscured 2.4 m telescope. These calibrations exist in the STScI Calibration Data Base, and are accessible with the STSDAS SYNPHOT package or with the XCAL software. The XCAL and SYNPHOT Users Guides should be consulted for further details. We include here a sufficient calibration for exposure planning. In Table 6.1 the dimensionless efficiency and the mean wavelength for each filter are tabulated together with the effective width, the equivalent Gaussian dimensionless width, the maximum transmission, the derivative of the mean wavelength with respect to spectral index, the pivot wavelength, average wavelength, and wavelength of maximum transmission. The parameters are defined as follows. The dimensionless efficiency is integral Q(lambda)T(lambda)dlambda/lambda The mean wavelength is defined in Schneider, Gunn, and Hoessel Ap. J., 264, 337 (1983) integral Q(lambda)T(lambda)log_e(lambda)dlambda/lambda lambda-bar = exp ------------------------------------------------------ integral Q(lambda)T(lambda)dlambda/lambda This rather unconventional definition has the property that the correspondingly defined mean frequency is just c/lambda-bar. It is in some sense halfway between the conventional frequency mean and the wavelength mean. The pivot wavelength is defined as integral Q(lambda)T(lambda)lambda d lambda lambda_p = ------------------------------------------ squared integral Q(lambda)T(lambda)dlambda/lambda The effective dimensionless Gaussian width is defined implicitly by sigma^2 = integral from lambda-bar(1+5sigma) to lambda-bar(1-5sigma) of Q(lambda)T(lambda)[log_e(lambda/lambda-bar]^(1/2) divided by integral from lambda-bar(1+5sigma) to lambda-bar(1-5sigma) of Q(lambda)T(lambda)(dlambda/lambda) where the integration limits eliminate unrealistic contributions from imperfect blocking far from the bandpass. The effective width of the bandpass is delta lambda-bar = 2[2log_e2]^(1/2) sigma lambda-bar The final two columns in Table 6.1 are defined as follows. In the next-to-last column m_e/sec is the zero-point magnitude for 1 e- s-1 (with AB_v=0). The final column gives t_wfsky, which is the exposure time (in seconds) needed to make the sky noise equal to 5 e- RMS (i.e. ~read noise) in the WFC for a sky level of V=23.3 mag arcsec^-2. Table 6.1: System Efficiencies and Zeropoints. See Section 6.1 for definitions. Table 6.2: AB_v as a Function of Wavelength. AB_v is defined as a color-dependent correction from V magnitude to AB magnitude at frequency v. Wavelength (A) runs along the top; spectral classes run down the leftmost column. The second column contains B-V See Section 6.3. 6.2 On-Line Exposure Time Calculator We note that most of the calculations below are incorporated in the on-line WFPC2 Exposure Time Calculator (ETC) program, which is available on the WFPC2 WWW pages at http://www.stsci.edu/ftp/instrument_news/WFPC2/wfpc2_top.html on the software tools page. To use this program, one merely fills out an "HTML" form giving the target information (e.g. magnitude and color), camera configuration (PC or WFC, desired gain setting, and filter), and either the exposure time or the desired signal-to-noise ratio. As of this writing there are separate HTML forms for point sources, extended sources, point source with background light, and extended targets with background light. After filling out the form the user then clicks on "calculate" and the program returns the resulting signal-to-noise ratio if the exposure time was specified, or vice versa. Examples of completed HTML forms and results are shown in "Sample SNR Calculations" on page 123. Note that clicking on any colored text on the HTML form will give a description of that item. The ETC program handles sources with stellar spectra, power law sources, and emission line sources; point sources and extended sources; and sources superposed on a diffuse stellar background. The program also returns advice on CR-SPLITing, use of CLOCKS=YES, and warnings about saturation, if appropriate. The program currently (Version 2.0) assumes stellar data will be analyzed by PSF fitting. Results are typically accurate to a few percent. While observers should familiarize themselves with the material below, most will find the ETC program faster and easier to use for actual calculations. The ETC program will also be updated to reflect any changes in instrument performance, so observers can be assured of up-to-the-minute information. 6.3 Target Count Rates We now consider estimation of count rates for objects with stellar, power law, and emission line spectra. 6.3.1 Count Rates for Stellar Sources To estimate the number of electrons collected from a point source of apparent visual magnitude V, one can use the equation: (6.1) N = 2.5 x 10^11 t [integral Q(lambda)T(lambda) dlambda/lambda] x 10^(-0.4(V+AB_v)) where t is the exposure time in seconds, the QT integral is given in Table 6.1, and AB_v is given in Table 6.2 as a function of spectral type and wavelength for some example spectral energy distributions. The quantity AB_v is a color-dependent correction from V magnitude to AB magnitude at frequency v. The AB magnitude system is defined as (Oke and Gunn 1983) AB = V + AB_v = -2.5 logf_v - 48.60 where F_v is the flux in erg cm^-2 s^-1 Hz^-1. Equation 6.1 may be trivially rewritten to give the count rate R_object in units of e- s^-1 pixel^-1 for a target with a stellar spectrum as: (6.2) R_object = 2.5 x 10^11[integral Q(lambda)T(lambda)dlambda/lambda] x 10^(-0.4(V + AB_v)) 6.3.2 Count Rates for Power Law Sources If one knows the spectral index a (which is zero for a source with a flat continuum), V+AB_v can also be calculated as the monochromatic Oke system magnitude at the corrected mean wavelength of the filter: V + AB_v = -2.5log_10(S_v[lambda-bar + alpha(dlambdabar/dlambda]) - 48.6 where S_v is the flux in ergs cm^-2 s^-1 Hz^-1 as in Oke and Gunn, Ap. J., 266, 713 (1983) at the effective mean wavelength of the filter lambda-bar + alpha(dlambda-bar/dlambda). It can be shown that dlambda-bar/dalpha = lambda-bar sigma^2 if the integrands are weighted by a source with spectral index (alpha in the definition of alpha. See also Koornneef, J., et al. "Synthetic Photometry and the Calibration of the Hubble Space Telescope" in Highlights of Astronomy (7, 833, J.-P. Swings Ed (1983). Combining the above equations gives (6.3) R_object = 6.9 x 10^30 [integral Q(lambda)T(lambda)dlambda/lambda] x S_v(lambda-bar + alpha dlambda-bar/dalpha) 6.3.3 Count Rates for Emission Line Sources The count rate in units of e- s^-1 for a monochromatic emission line is given by (6.4) R_object = 2.3 x 10^12. (QT) x F x lambda where F is the emission line flux in units of ergs cm^-2 s^-1 , and lambda is the wavelength of the line in Angstroms. The quantity QT is the (system + filter) quantum efficiency at the wavelength of the line, which can be determined from inspection of the figures in section A.2, "Passbands including the System Response", on page 211. For lines near the maxima of the filter transmission curves, it should be sufficient to use QT. from Table 6.1. Note that the integrated filter efficiency is not relevant for the signal calculation. In cases where the width of the line approaches that of the filter, it will be necessary to convolve the line shape and filter bandpass using either the SYNPHOT or XCAL programs. For example, H_alpha emission at 6563A, with total source flux F=10^-16 erg s^-1 cm^-2, observed through the F656N filter (total system throughput T=0.104 from "F656N, F658N, F673N, F675W, F702W, F785LP" on page 217), will produce a target count rate R_object@ = 155 e- s^-1 integrated over the entire source. 6.4 Sky Background The sky background can contribute significant Poisson noise in broad and medium band filters, and must be taken into account during noise calculations. The actual sky brightness depends on the heliocentric ecliptic coordinates (latitude and longitude) in a manner summarized in Table 6.3. The appropriate AB_v can be taken from Table 6.2. To convert mag arcsec^-2 to mag pixel^-1 one needs to add 5 magnitudes (WFC) or 6.7 magnitudes (PC1). These values are actually lower limits on the effective sky-brightness that will be seen, because light from the bright Earth limb can scatter into the aperture. If your observations are sky background limited, and signal-to-noise is a driver, consider the use of the special requirement LOW-SKY as described in the Call for Proposals or the Phase II Proposal Instructions. LOW-SKY has two effects: * It causes the observation to be scheduled at the time of year when the zodiacal background light is within 30 percent of the minimum possible background value for the target, and * It requires that the observation be made when the bright Earth limb is more than 40 degrees from the OTA axis, which greatly reduces scattered light. For many targets LOW-SKY will have minimal impact on the observing efficiency. Note, however, that targets in the Continuous Viewing Zone (CVZ) cannot be observed if LOW-SKY is specified. See "Observing Faint Targets" on page 143 for more information. Table 6.3: Sky Brightness (V mag arcsec^-2) as a Function of Heliocentric Ecliptic Latitude and Longitude. "SA" denotes that the target is unobservable due to solar avoidance. Heliocentric Ecliptic Latitude (degrees) Ecliptic Longitude 0 15 30 45 60 75 90 (degrees) ------------------------------------------------------------------------------ 180 22.1 22.4 22.7 23.0 23.2 23.4 23.3 165 22.3 22.5 22.8 23.0 23.2 23.4 23.3 150 22.4 22.6 22.9 23.1 23.3 23.4 23.3 130 22.4 22.6 22.9 23.2 23.3 23.4 23.3 120 22.4 22.6 22.9 23.2 23.3 23.3 23.3 105 22.2 22.5 22.9 23.1 23.3 23.3 23.3 90 22.0 22.3 22.7 23.0 23.2 23.3 23.3 75 21.7 22.2 22.6 22.9 23.1 23.2 23.3 60 21.3 21.9 22.4 22.7 23.0 23.2 23.3 45 SA SA 22.1 22.5 22.9 23.1 23.3 30 SA SA SA 22.3 22.7 23.1 23.3 15 SA SA SA SA 22.6 23.0 23.3 0 SA SA SA SA 22.6 23.0 23.3 ------------------------------------------------------------------------------ Another option for reducing the sky brightness, is the special requirement SHADOW, which forces the observation to be made when HST is in the Earth's shadow. This usually has a large negative impact on the observing efficiency, and is recommended only when attempting to avoid geocoronal lines when observing far-UV emission lines (e.g. Ly alpha and OI 1304A). Moreover, it does not attempt to minimize zodiacal emission, which dominates at visible wavelengths. Table 6.4 shows approximate sky count rates for the WFC and PC1 for filters with significant sky count rates. An average sky brightness of V=22.9 mag arcsec^-2 is assumed. Filters not listed in the table have sky count rates below that of the dark current, so the sky contribution will generally be unimportant. Values for other filters or sky brightnesses can be computed from Table 6.2, Table 6.1, Table 6.3, and Equation 6.2. Table 6.4: Sky Count Rate per Pixel (Psky). An average sky brightness of V = 22.9 mag arcsec^-2 is assumed. Filters not listed have sky rate significantly below the dark current. Filter Sky Count Rate (P_sky) (e- s^-1 pixel^-1) WFC PC1 ---------------------------------------- F336W 0.0009 0.0002 F38OW 0.005 0.001 F439W 0.005 0.0011 F45OW 0.018 0.004 F467M 0.003 0.0006 F547M 0.021 0.0045 F555W 0.052 0.010 F588N 0.002 0.0006 F569W 0.040 0.0081 F606W 0.090 0.020 F622W 0.060 0.012 F673N 0.002 0.0006 F675W 0.056 0.012 F702W 0.082 0.0016 F785LP 0.024 0.0050 F791W 0.048 0.010 F814W 0.054 0.011 F85OLP 0.012 0.0024 ---------------------------------------- 6.5 Signal-to-Noise Ratio Estimation The signal-to-noise ratio (SNR) for a point source depends on both the Poisson noise of the object, and on noises associated with the background. Sources of background noise include "read noise" of the CCDS, and Poisson noise in the dark current, sky background, and any smooth galaxy light superposed on the target. The SNR obtained for photometry of a point source will depend on the analysis technique used. The optimum SNR will be obtained when the pixels of the point source PSF are weighted in proportion to their expected intensity by PSF fitting. Aperture photometry will tend to give lower SNR, especially for sources where the background is important, but nonetheless is widely used. We now consider both methods. 6.5.1 Point Sources -- PSF Fitting In the bright target limit, Poisson noise sets the SNR and SNR = (S)^(1/2) = (R_object x t)^(1/2) where S is the number of detected photons, and R_object is given by the above equations 6.2 through 6.4, and t is the exposure time. In the background limited case (e.g. read noise, dark current, or sky noise limited) the SNR is a function not only of the expected number of detected photons S from the source but also of the average effective background count rate in each pixel B, the point spread function (PSF)_i,j, and the weights used to average the signal in the pixels affected by the source. It is easy to show that the signal-to-noise ratio for optimal weights (which are proportional to the point spread function) is given by: (6.5) SNR = S/square root of B times (SUM [PSF_i,j^2])^(1/2) = S/square root of B times sharpness^(1/2) Where sharpness is effectively the reciprocal of the number of pixels contributing background noise. The summation is tabulated for a few representative cases in Table 6.5. To estimate the signal-to-noise, multiply the signal-to-noise obtained, assuming all the flux is in one pixel, by the square root of the value in the Table. Table 6.5: Sharpness as a Function of Wavelength, Camera, and Location of the Star Center with Respect to the Pixel Grid. The "Obs" columns represent the values for the real OTA, WFPC2 optics, and CCD MTF function. The "Diff column represents values for the theoretical diff raction limit with perfect optics and detectors. Target location refers to both the camera used (PC or WFC), and the location of the star center on the pixel grid. Target Location 2000 A 4000 A 6000 A 8000 A Obs. Diff. Obs. Diff. Obs. Diff. Obs. Diff. ------------------------------------------------------------------------------ PC Pixel Center 0.084 0.409 0.095 0.259 0.066 0.115 0.046 0.073 PC Pixel Comer 0.063 0.186 0.065 0.107 0.054 0.072 0.045 0.068 WFC Pixel Center 0.120 0.745 0.145 0.482 0.128 0.318 0.124 0.285 WFC Pixel Comer 0.102 0.228 0.105 0.193 0.098 0.178 0.081 0.126 ------------------------------------------------------------------------------ We note that PSF fitting is equivalent to convolving the image with the PSF, and then measuring the peak counts for stellar objects. Also, the location of the star on the pixel grid will be impossible to know in advance of the observation (i.e. pixel center vs. pixel corner in Table 6.5). In general, the lower "pixel corner" values should be used, so as to insure adequate SNR. The average effective background counts per exposure and per pixel can be expanded to include various sources: B = readnoise^2 + P_dark x (t + 46) + P_sky * t + P_background x t where terms include the read out noise of the CCD (readnoise), the dark current (P_dark), sky background count rate (P_sky), and the count rate of any diffuse background light from astrophysical sources (P_background). Herein we will use "P" to represent count rates per pixel, and "R" to represent the total counts for an object. The exposure time is represented by t. For example, Table 2.2 lists the faintest V magnitude star, V=28.14, measurable with a signal-to-noise ratio of 3 in a 3000s integration in F569W in the Wide Field Cameras. The calculation to check this goes as follows. The efficiency of the filter is 0.02139 from Table 6.1. The sky background in each pixel is 23.3+5=28.3, assuming an ecliptic latitude of 90 degrees from Table 6.3, and the pixel area correction for the WFC given in that section. The total sky background collected per pixel in 3000 seconds is given by Equation 6.1 as 76.8 electrons. Note that the AB_v color correction required for the sky in the wavelength range of the filter is 0.0 from Table 6.2. From Table 4.3, the read noise is 5.2 electrons. From Table 4.2, the median dark current at -88 'C is 0.004. Therefore the total dark current (on which there will be shot noise) is only 12 electrons. The equivalent background per pixel is then given as B=76.8+5.2^2+12=116. The total number of detected electrons from a star with V=28.14 is S=89 electrons, again using Equation 6.1. The expected peak count is 27 detected electrons using Table 5.4, which is much less than B, requiring the use of Equation 6.5 for the background limited case. The sharpness for the WF camera in the best case, when the star is centered on a pixel, is given in Table 6.5 as 0.128. Then Equation 6.5 above gives the signal-to-noise as 3.0: SNR=(89/square root 116) x (square root 0.128) = 3.0 If, instead, the peakcount rate comes out much greater than the background, the observation is photon noise limited, and the signal-to-noise should be computed as the square root of the signal S in electrons. In principle, one should also include contributions in the signal-to-noise for flat fielding uncertainties, noise in the bias and dark calibration files, and quantization noise. Flat fielding errors will be of order 1 percent, and will limit SNR in the large-signal limit. Noise in the bias and dark calibration files will be unimportant in most pixels, although these could become important if many (> 10) non-dithered frames of the same field are combined. Quantization noise can be estimated as gain/square root 12 (i.e., square root 4.l in the 7 e- DN^-1 channel, and square root 16.3 in the 14 e- DN^-1 channel). In nearly all situations it can be ignored. In the weak signal case, the quantization noise is effectively included in the read noise values given throughout this Handbook; in the strong signal case it is very small compared to the Poisson noise and can be ignored. A generalized equation for estimating point source signal-to-noise ratio per exposure is given below (Equation 6.6). It is exact in both the bright and faint object limits, and is a reasonable approximation to the intermediate case. P_background represents any generalized source of diffuse background light (e.g. galaxy on which target is superposed). Table 6.6 gives rough values for some of the parameters, along with references for more accurate values. (6.6) Letting: P = R_object x t Q = readnoise^2 + P_dark x (t + 46) + P_sky x t + P_background x t SNR = P/(square root[P + (Q/sharpness)]) Note that in this formulation, sharpness^-1 is the equivalent number of pixels the weighted signal is integrated over. In the event that multiple exposures are taken (e.g. to remove cosmic rays), the signal-to-noise ratio for the final averaged image is approximately given by: SNR_total = SNR x square root N where N is the number of images averaged. Table 6.6: Parameters for Point Source SNR Estimation - PSF Fitting Parameter Description Units Approx. Value Better Value ------------------------------------------------------------------------------ R_object object count e- s^-1 Equation 6.1, rate 6.2, or 6.3 P_dark dark count e- s^-1 pixel^-1 0.004 Table 4.2 rate P_sky sky count e- s^-1 pixel^-1 Table 6.4 Tables 6.2, rate 6.1, 6.3; Eqn 6.1 P_background count rate e- s^-1 pixel^-1 Tables 6.2, from back- 6.1; Eqn 6.1 ground light (if any) read noise e- ATD-GAIN=7 use 5.3^a Table 4.3 ATD-GAIN=15 use 7.5 sharpness WFC use 0.11 Table 6.5 PC1 use 0.06 t exposure time s ------------------------------------------------------------------------------ a. ATD-GAIN defaults to 7 unless otherwise specified on Phase 2 proposal. 6.5.2 Point Sources -- Aperture Photometry When aperture photometry is used, one must consider the fraction of the object counts encircled by the aperture, as well the background noise in the aperture. In the bright target limit the SNR is given by SNR = (S x f(r))^(1/2) = (R_object x f(r) x t)^(1/2) where S is the number of detected photons, f(r) is the fraction of the total counts encircled by the aperture with radius r, and R_object is target count rate. Representative values of f(r) are given in Table 6.7; values for other aperture sizes and filters can be estimated from Figure 5.3 on page 88, or Figure 5.4 on page 89. In the faint target limit the noise contributed by background counts determines the SNR SNR = (S x f(r))/(square root[B x pie x r^2]) where B represents the effective background counts per pixel, and r is the aperture radius in pixels. In the generalized case the SNR per exposure for aperture photometry is given approximately by: (6.7) Letting: P = (f(r) x R_object x t) Q = (readnoise^2 + P_dark x (t + 46) + P_sky x t + P_background x t) SNR = P/(square root[P + Q x pie r^2]) where the parameters are summarized in Table 6.8. Table 6.7: Encircled Energy for Representative Filters. Encircled energy values are normalized to unity at large radius. CCD Aperture Radius Encircled Energy f(r) (r) F218W F555W F814W ------------------------------------------------------------------------------ PC1 0.1" 0.60 0.67 0.53 0.2" 0.73 0.85 0.78 0.5" 0.84 0.96 0.87 1.0" 0.92 1.00 0.92 WF3 0.1" 0.40 0.46 0.44 0.2" 0.69 0.76 0.74 0.5" 0.85 0.90 0.91 1.0" 0.94 0.94 0.96 ------------------------------------------------------------------------------ Table 6.8: Parameters for Point Source SNR Estimation - Aperture Photometry Parameter Description Units Approx. Value Better Value ------------------------------------------------------------------------------ R_object object count e- s^-1 Equation 6.1, rate 6.2, or 6.3 P_dark dark count e- s^-1 pixel^-1 0.004 Table 4.2 rate P_sky sky count e- s^-1 pixel^-1 Table 6.4 Tables 6.2, rate 6.1, 6.3; Eqn 6.1 P_background count rate e- s^-1 pixel^-1 Tables 6.2, from back- 6.1; Eqn 6.1 ground light (if any) readnoise e- ATD-GAIN=7 use 5.3^a Table 4.3 ATD-GAIN=15 use 7.5 f(r) encircled Table 6.7 Figure 5.3 energy or 5.4 r aperture pixels radius t exposure time s ------------------------------------------------------------------------------ a.ATD-GAIN defaults to 7 unless otherwise specified on Phase 2 proposal. 6.5.3 Extended Sources The calculations for extended sources are nearly identical to those for point sources. The easiest procedure is to compute the SNR per detector pixel, and then adjust this value if the total SNR is required for an area encompassing many pixels. In general, one will have the target magnitude or flux per square arcsecond. To compute the flux per pixel for the PC one merely multiplies the flux per square arcsecond by 0.00207, or instead, adds the value 6.7 to the magnitude per square arcsecond to get the necessary magnitude per PC pixel. For the WFC, one either multiplies the flux per square arcsecond by 0.00993, or adds 5.0 to the magnitude per square arcsecond. Equations 6.2, 6.3, and 6.4 can be rewritten including these factors as below. PC Camera For the PC camera, sources with stellar spectra, and V surface brightness per square arcsecond sigma_v we have a count rate in e- s^-1 pixel^-1 of (6.8) P_object = 2.5 x 10^11 x [integral(Q(lambda)T(lambda)dlambda/lambda] x 10^(-0.4(sigma_v + AB_v + 6.7)) For power law sources where B_v is the target flux in units of ergs cm^-2 s^-1 Hz^-1 arcsec we have (6.9) P_object = 1.4 x 10^28 x [integral(Q(lambda)T(lambda)dlambda/lambda] x B_v(lambda-bar + alpha(dlambda-bar/dalpha)) And finally for emission line sources where I_v is the flux in ergs cm^-2 s^-1 arcsec^-2 we have (6.10) P_object = 4.8 x 10^9 x (QT) x I_v x lambda where the emission line wavelength lambda is in Angstroms. WFC Camera For the WFC camera and stellar sources with V surface brightness per square arcsecond sigma_v we have a count rate in e- s^-1 pixel^-1 of (6.11) P_object = 2.5 x 10^11 x [integral(Q(lambda)T(lambda)dlambda/lambda] x 10^(-0.4(sigma_v + AB_v + 5)) For power law sources where B_v is the target flux in units of ergs cm^-2 s^-1 Hz^-1 arcsec^-2 we have (6.12) P_object = 6.9 x 10^28 x [integral(Q(lambda)T(lambda)dlambda/lambda] x B_v(lambda-bar + alpha(dlambda-bar/dalpha)) And finally for emission line sources where I_v is the flux in ergs cm^-2 s^-1 arcsec^-2 we have (6.13) P_object = 2.3 x 10^10 x (QT) x I_v x lambda where the emission line wavelength percent is in Angstroms. SNR The generalized SNR per pixel per exposure for an extended source is then obtained simply by setting the sharpness to unity in equation 6.5: (6.14) SNR = P_object x t divided by: square root[P_object x t + (readnoise^2 + P_dark x (t + 46) + P_sky x t + P_background x t)] Table 6.9: Parameters for Extended Source SNR Estimation. Parameter Description Units Approx. Value Better Value ------------------------------------------------------------------------------ R_object object count e- s^-1 pixel^-1 Equation 6.9 to rate 6.12 P_dark dark count e- s^-1 pixel^-1 0.004 Table 4.2 rate P_sky sky count e- s^-1 pixel^-1 Table 6.4 Tables 6.2, rate 6.1, 6.3; Eqn 6.7 (PC) or 6.10 (WFC) P_background count rate e- s^-1 pixel^-1 Tables 6.2, from back- 6.1; ground light Eqn 6.7 (PC) or (if any) 6.10 (WFC) readnoise e- ATD-GAIN=7 use 5.3^a Table 4.3 ATD-GAIN=15 use 7.5 t exposure time s ------------------------------------------------------------------------------ a. Default value is ATD-GAIN=7. Since many observations of extended sources are for galaxies in broad-band filters, a few rules of thumb can be useful. Saturation is seldom a concern, except in very bright spots such as the inner core of ellipticals and of some bulges. Count rates for spiral galaxies range typically from 2 to 0.01 e- pixel^-1 s^-1 (and lower) for filters such as F555W, F606W, F702W, and F814W; the lower end of the range corresponds roughly to the de Vaucouleurs D_25. Count rates are significantly lower in blue and UV filters. Spiral structure can typically be traced reasonably well with total exposures of 3000 seconds or longer in the above filters. For galaxies of very small angular size at red-shifts of cosmological interest, the image may cover a small number of pixels; thus the detection of such objects follows rules similar to those of point sources. However, the fraction of light falling in the central pixel is smaller for most galaxies than it is for true point sources. The approximate magnitude difference between the light falling in the central pixel and the entire galaxy is plotted in Figure 6.1 for a typical giant elliptical galaxy, as a function of redshift. For other types of galaxies, a morphological term can be added to the values (for example, 0.6 magnitudes for lenticulars, 0.7 for S, 0.8 for Sab, 0.9 for Sbc, 1.2 for Scd, and 1.8 for Irr). These values must be increased by an additional 1.7 magnitudes for the PC. Figure 6.1: Giant Elliptical Galaxy 6.6 Exposure Time Estimation In many instances one desires a certain SNR, and wishes to solve for the corresponding exposure time. Given the SNR, Equations 6.6, 6.7, or 6.14 can be solved for the exposure time, t. Since there are time-dependent and time-independent noise sources, quadratic equations are obtained. For example, we may solve equation 6.6 for the point source exposure time: t = 1/2Y x (b + square root[b^2 + 4aY]) where the term A contains the time-independent noise sources a = (readnoise^2 + 46P_dark)/sharpness and B contains the time-dependent noise sources b = [(P_dark + P_sky + P_background)/sharpness] + R_object and y = (R_object/SNR)^2 Equations for aperture photometry (6.7) and extended sources (6.12) can be solved with similar results. Parameters are as described in Tables 6.6, 6.8, and 6.9. We again note that the on-line WFPC2 Exposure Time Calculator program provides an easy method for these calculations. 6.7 Sample SNR Calculations Below we give further examples of SNR calculations. The Appendix also gives tables of SNR values for a wide range of representative cases. 6.7.1 Point Sources Simple Star, Manual Calculation, PSF Fitting We begin with the simple example of V=20 star of spectral class G0. We want to observe with the PC using filter F555W. The star is somewhere near the ecliptic pole. We want to know the SNR for a 1200s CR-SPLIT exposure. Default ATD-GAIN=7 is used. We plan to use PSF fitting to analyze the data. First we estimate the count rate for our target. Consulting Equation 6.2, Table 6.1, and Table 6.2 we have: R_object = 2.5 x 10^11 x [integral(Q(lambda)T(lambda)dlambda/lambda)] x 10^(-0.4(V+AB_v)) = 2.5 x 10^11 x [0.0271] x 10^(-0.4(20 + 0.02) = 66 in units of e- s^-1. Next we fill out Equation 6.6. To keep things simple we just use values from Table 6.6, and get the sky count rate from Table 6.4. There is no background light (i.e. no superposed galaxy), so P_background=0. The exposure time t=600 for each exposure of the CR-SPLIT: Letting: J = (R_object x t) K = (readnoise^2 + P_dark x (t+46) + P_sky x t + P_background x t) SNR = J/(square root[(J + K)/sharpness]) = (66x600)/square root[66x600+((53.3)^2+.004x(600+46)+(.01x600)+(0x600))/.06] = 39600/square root[39600+611] = 196 The SNR for the total 1200s exposure, i.e. both halves of the CR-SPLIT, would simply be: SNR_total = SNR x square root N = 197 x square root 2 = 279 At these high SNR levels, it is likely that flat fielding would limit the photometric accuracy, rather than the noise. If we have a look at the terms in the SNR equation, we can see that the Poisson noise dominates; the term containing the sharpness and background noise sources is unimportant. Just for fun, let us see what happens if we keep everything the same, but give the target V=25. Now we have R_object=0.66 e- s^-1 , and: SNR = (.66x600)/square root[.66x600+((53.3)^2+.004x(600+46)+(.01x600)+(0x600))/.06] = 396/square root[396+611] = 12.5 We see that now the term with the background noise (in particular, the read noise) limits the SNR. For the full 1200s exposure the SNR_total=17.6. Simple Star, Manual Calculation, Aperture Photometry What if we now want to observe this same V=25 star, but we plan to reduce the data by measuring counts in a 0.5" radius aperture? We now use Equation 6.7 Letting: J = (f(r) x R_object x t) K = [readnoise^2 + P_dark x (t + 46) + P_background x t] SNR = J/square root[ J + K x pie r^2] (.96x.66x600) = --------------------------------------------------------------------------- square root[(.96x.66x600)+((5.3)^2+.004(600+46)+.01x600+0x600xpie(11.6)^2] = 380/square root[380 + 15500] = 3.0 Apparently using aperture photometry with a 0.5" radius aperture reduces the SNR by a factor ~4 as compared to PSF fitting, for this background limited case. Simple Star, SNR Tables, PSF Fitting We now repeat the first calculation above for the V=20 star using the SNR tables in Appendix B. We look up the GO spectral class and F555W filter (5500,A) in Table B.1, and obtain AB_v=0.02. For the V=20 star, we thus have V+AB_v=20.02. We look at Figure B.10 on page 234 and find this value on the horizontal axis. We locate exposure time 600s (one-half of the total 1200s CR-SPLIT exposure), and find SNR ~200. For the total 1200s exposure the SNR would be 200(square root[2] = 280. Simple Star, On-Line Calculator, PSF Fitting The above calculation for a V=20 GO star may also be performed using the WFPC2 Exposure Time Calculator program, which is available on the WFPC2 WWW pages at: http://www.stsci.edu/ftp/instrument_news/WFPC2/wfpc2_top.html To use this program, access the above address with NCSA Mosaic, Netscape, or a similar program. Once in the WFPC2 area, select the "Software Tools" page, and then the "ETC" page. For the first example above, choose the "Point Source" form and complete it as shown in Figure 6.2 for the 600s sub-exposure. Then clicks the "calculate" button and after a few seconds the result is displayed (Figure 6.3). The answer, SNR=198, is comparable to that obtained by the manual calculation above for the 600s sub-exposure (SNR=197). Alternatively, one can input the total exposure time (1200s), and then use the result farther down the output page for "No. Sub-Exposures = 2" (see Figure 6.4), thereby obtaining SNR=277 for the total 1200s CR-SPLIT exposure. Figure 6.2: Sample Fill-out Form for WFPC2 On-Line Exposure Time Calculator. Figure 6.3: Sample Results from WFPC2 On-Line Exposure Time Calculator. Figure 6.4: Sample Results on CR-SPLITing from WFPC2 On-Line Exposure Time Calculator Results Page. Star Superposed on Galaxy, Manual Calculation We now consider a B=25 point source of spectral class B0, which is superposed on an elliptical galaxy with sigma_v=22 mag arcsecond. We want to compute the SNR obtained from a one-orbit (40 min.) non-CR-SPLIT observation in filter F300W on the WFC. PSF fitting will be used for the photometry. We begin by computing the total count rate for the target. Using Table 6.2 we see that this target will have V=25.31. From Table 6.1 we obtain the filter efficiency and mean wavelength. Interpolating by mean wavelength in Table 6.2 we obtain AB_v=-0.83 for the B0 star. Using Equation 6.2 we have: R_object = 2.5 x 10^11 x (integral[Q(lambda)T(lambda)dlambda/lambda]) x 10^(-0.4(V+AB_v)) = 2.5 x 10^11 x [0.00519] x 10^(-0.4(25.31 - 0.83)) = 0.21 in units of e- s^-1. Next we consider the background light from the superposed galaxy. We set sigma_v=22 mag arcsecond^-2 in Equation 6.11, and AB_v=3.63 for a gE galaxy at lambda=3000A (filter F300W) from Table 6.2. Hence the count rate per pixel due to the background light is: P_background = 2.5x10^11x(integral[Q(lambda)T(lambda)dlambda/lambda]x10^(-.4(sigma_v+AB_v+5)) = 2.5 x 10^11 x [0.00519] x 10^(-0.4(22 + 3.63 + 5)) = 0.00073 For the sky background, we note that Table 6.4 has no entry for F300W, so that the sky must be unimportant. If we wanted to calculate it anyway, as a check, we would use Table 6.3 for the sky brightness, Table 6.2 for the sky's AB_v, and again Equation 6.11. We will assume the target is near the ecliptic pole. P_sky = 2.5x10^11x(integral[Q(lambda)T(lambda)dlambda/lambda]x10^(-.4(sigma_v+AB_v+5)) = 2.5 x 10^11 x [0.00519] x 10^(-0.4(23.3 + 3.12 + 5)) = 0.00035 For the sharpness function we will use "pixel corner" values (least optimistic choice) from Table 6.5. Using read noise and dark current from Table 6.6, and Equation 6.6 for point source SNR: Letting: J = (R_object x t) K = [readnoise^2 + P_dark x (t + 46) + P_sky x t + P_background x t] SNR = J/square root[ J + K/sharpness] = 0.21 x 2400 ------------------------------------------------------------------------------ squareroot[.21x2400+((5.3)^2+.004x(2400+46)+(.00035x2400)+(.00073x2400))/.103] for this single exposure. The SNR for multiple 40 min. exposures would be simply 16.8(N^(1/2)), where N is the number of exposures. Star Superposed on Galaxy, On-Line Calculator The above calculation could also be performed with the on-line WFPC2 Exposure Time Calculator. One would select the "Point source + stellar background" form, and complete it as in Figure 6.5, and then click on "calculate." Figure 6.6 shows some of the results. 6.7.2 Extended Sources In general, the signal-to-noise level for extended sources can be computed by comparing the expected signal, S, in each pixel, determined from Equations 6.8 through 6.13, to the noise N=(S+B)^(1/2), where B is the equivalent background, determined in a manner similar to that for point sources. Unlike for point sources, the calculation does not, in a first approximation, involve the sharpness of the point spread function. For example, let us consider the observation of a source with a V surface brightness of 24 mag arcsec- , assuming the F569W filter F569W, WFC camera, and sky background V=23.3 mag arcsec^-2. The signal-to-noise estimate goes as follows. The signal in each WFC pixel is 24.0+5.0 = 29.0 magnitude. By Equation 6.11, the total signal collected from the source in a 3000 second integration is S = 41.8 electrons, neglecting the small AB color correction. The sky signal per pixel is 76.8 electrons. The dark current is ~12 electrons. The total equivalent background is thus B = 76.8 +5 .3^2 + 12 = 116.9 electrons, larger than the Signal detected, thus the noise is background-dominated. The noise is N=(S+B)^(1/2) = 12.6 electrons, and the signal-to-noise per pixel expected in this case is 3.3. Similar calculations can be carried out for other filters; for observations in narrow-band filters and in the UV, the sky background signal will usually be unimportant. For very long observations of faint objects, other noise terms, such as flat field uncertainty, and errors in dark (and possibly bias) subtraction, must be considered more carefully. If the scale of features in the target is larger than one pixel, the signal-to-noise can sometimes be improved by smoothing the observed image or - if read noise is a significant contributor - by reading the image out in AREA mode (see section 2.8, "CCD Orientation and Readout", on page 3 1). 6.7.3 Emission Line Sources The signal-to-noise ratio calculation for point-like or extended emission-line sources is similar to that for continuum sources. However, the details of the calculation are different, because of the units used for the line flux, and because the flux is in a narrow line. The integrated filter efficiency is not relevant for the signal calculation; what matters is the total system throughput QT at the wavelength of the line, which can be determined from inspection of the Figures in section A.2, "Passbands including the System Response", on page 211. For lines near the center of the filter bandpasses the QT. values from Table 6.1 can be used. The total signal expected for a point source of line strength F, expressed in erg s^-1 cm^-2, is S=2.28x10^12 lambda t QT F, where t is the exposure time in seconds, and lambda the wavelength of the line in Angstroms. Thus, H. emission at 6563A, with flux F=10^-16 erg s^-1 cm^-2, observed for 1000 seconds through the F656N filter (total system throughput QT=0.104 from "F656N, F658N, F673N, F675W, F702W, F785LP" on page 217), will produce a total signal of S=155 electrons. The equivalent background per pixel is read-noise dominated: B=1+5.32^2+4=33, for a background noise of ~6 electrons. The total noise is dominated by photon noise from the signal itself, and the signal-to-noise ratio achieved in this observation is over 10. If the source is extended, the expected signal per arcsecond must be multiplied by the effective pixel area: 0.0099 arcsec^2 for the WF, 0.0021 for the PC. For a line flux of, say, F = 10-15 erg s^-1 cm^-2 arcsec^-2, this corresponds to 15 electrons in 1000 seconds for a WFC pixel. The noise is now dominated by the background, and the single-pixel signal-to-noise ratio is 15/(33 + 15)^(1/2) ~ 2.1. Figure 6.5: Point Source + Stellar Background Fill-out Form for WFPC2 On-Line Exposure Time Calculator. SNR is calculated for B=25 star (class B0) superposed on an elliptical galaxy (gE) with av=22. WFC is used with F300W. Figure 6.6: Sample Output from WFPC2 On-Line Exposure Time Calculator. Extended Line Emission Source, Manual Calculation We now consider a detailed example of a planetary nebula observed on the PC with the F502N filter. The nebula has a diameter of 5" and a total flux F=4x 10^-13 erg s^-1 cm^-2 in the [OIII] 5007A line. We want to estimate the SNR for an 1800s exposure, which will be CR-SPLIT. First we must estimate the flux per square arcsecond. Using the nebula diameter, the average brightness is I_v = 2.O x 10^-14 erg s^-1 cm^2 arcsec^-2. From the plots in section A.2, "Passbands including the System Response", on page 211, we see that QT=0.053. Using Equation 6.10 for the target count rate per pixel: P_object = 4.8 x 10^9 x (QT) x I_v x lambda = 4.8 x 10^9 x (0.053) x 2.0 x 10^9-14 x 5007 = 0.025 Next we estimate the SNR for each 900s sub-exposure using Equation 6.14 and Table 6.9. For this narrow filter the sky background can be ignored. We presume there is no background light from astrophysical sources: Letting: J = (P_object x t) K = [readnoise^2 + P_dark x (t + 46) + P_sky x t + P_background x t] SNR = J/square root[ J + K] = 0.025 x 900 ------------------------------------------------------------------ square root[0.025 x 900 + 5.3^2 + .0004 x (900 + 46) + 0 x 900 = 3.1 Hence SNR=3.1 per pixel for each 900s sub-exposure. The SNR per pixel for the total 1800s is SNR_total = SNR square root[N] = 3.1 x square root[2] = 4.3 The SNR for the entire nebula is this SNR per pixel times the square root of the number of pixels in the image, or ~420. In actuality, uncertainties in the photometric calibration and flat fields, would limit the SNR to ~100. Extended Line Emission Source, On-Line Calculator The above example could be calculated with the "Extended Source" form of the ETC program. The fill-out form would be completed as shown in Figure 6.7. We have selected "[OIII] 5007" on the emission line menu, and have left the redshift (z) set to zero. The PC and F502N filter are selected. Note we have entered the exposure time as 1800s. Scrolling down through the output page we find a table of SNR for various CR-SPLITings of the exposure (See Figure 6.8). "No. Sub-Exposures = 2" gives the answer we want, SNR=4.5 per pixel. Figure 6.7: Extended Source Form for WFPC2 On-Line Exposure Time Calculator. Here target is galactic [OIII] 5007 line emission source and is observed on PC with filter F502N. SNR is computed for 1800s exposure. Figure 6.8: Sample Output from WFPC2 On-Line Exposure Time Calculator. Line Emission Point Source w/ LRF, Manual Calculation In this example we consider an unresolved source of H_alpha emission in a galaxy at redshift z=0.22 with flux F=1.5x10^-16 erg s^-1 cm^-2. We want the SNR for a 2400s exposure without CR-SPLITing. Since the redshift is significant, we cannot observe with the F656N filter. Instead we will use the Linear Ramp Filter (LRF). The observed wavelength will be 8007A. From Table 3.7 on page 48 we see that this will be observed using the FR868N filter on CCD WF3. Combining the LRF transmission from Figure 3.2 and the "VVFPC2 + OTA System Throughput" from Figure 2.4 we estimate QT=0.054. We compute the count rate using Equation 6.4: To estimate the SNR R_object = 2.3 x 10^12 x (QT) x F x lambda = 2.3 x 10^12 x (0.054) x (1.5 x 10^-16) - 8007 = 0.15 we use Equation 6.6, which assumes that PSF fitting will be used to analyze the image. Since the filter is narrow, we will ignore the sky emission. We use Table 6.6 for the WFC sharpness and also the read noise. Letting: J = (P_object x t) K = [readnoise^2 + P_dark x (t + 46) + P_sky x t + P_background x t] SNR = J/square root[J + (K/sharpness)] = 0.15 x 2400 ------------------------------------------------------------------------------ square root[.15 x 2400 + (5.3^2 + 0.004(2400+46) + 0 x 2400 + 0 x 2400)/.11)] = 360/square root[360 + 344] = 14 which is for an un-split 2400s exposure. The Poisson noise and background noises contribute nearly equally. For three such exposures over three orbits SNR_total = SNRV square root[N] = 14 square root[3] = 23. Line Emission Point Source w/ LRF, On-Line Calculator The above calculation can be performed using the ETC program. The "Point Source" form is used. "Emission Line" source and the "H 6563" line are selected; the redshift is set to 0.22. The program will automatically choose between PC and WFC, depending on the LRF setting. The least optimistic case of placing the object on a "pixel corner" is selected. The filter "LRF" is selected from the filter menu, and 8007A is given for the central wavelength. The exposure time is specified as 2400s. (See Figure 6.9 for example of completed form.) The result is SNR=13.1 for the un-split 2400s exposure (Figure 6.10), which is comparable to the manual calculation of SNR=14. Figure 6.9: Point Source Form for WFPC2 On-Line Exposure Time Calculator. Target is unresolved galaxy (z=0.22) nucleus with Halpha line emission which is observed with LRF. SNR is computed for 2400s exposure. Figure 6.10: Sample Output for WFPC2 On-Line Exposure Time Calculator. 6.8 Red Leaks in UV Filters The presence of significant red leaks in the UV filters, together with the much greater sensitivity and wavelength coverage in the red part of the spectrum, can make UV observation and calibration difficult. Observers must sometimes be prepared to take additional frames at red wavelengths, in order to estimate the contribution of red leak to the UV counts. The counts contributed by red leak can be a significant noise source, and must also be taken into account during SNR and exposure time estimation. See section 3.7 on page 57 for detailed information. Note that the SYNPHOT synthetic photometry package can be used to estimate counts due to red leak for particular filter / target combinations. 6.9 Time Dependence of UV Response The UV throughput of the WFPC2 degrades in a predictable way after each monthly decontamination. The photometric calibration given in Section 6.1 is applicable at the start of each cycle, and measurements taken at other times must be corrected to account for the change in sensitivity since the last decontamination. In addition, a long-term change in sensitivity is present for the F160BW and F170W filter observations on the PC, and may be present to a lesser degree at other wavelengths. Figure 6.11 shows the photometric monitoring data for the standard star GRW+70D5824 (a white dwarf classified DA3; B-V = -0.09) in the WF3 and PC1 for the set of filters which are routinely monitored. The dotted vertical lines mark the dates of the decontaminations. The solid vertical lines labeled "-88" mark the date the operating temperature was changed from -76 degrees C to -88 degrees C (April 24, 1994). The following information is for the period following the cooldown date. Figure 6.11 shows that the effect of contamination on the F675W and F814W filter observations is essentially negligible, but it rises toward the UV until it reaches values of 30 percent to 40 percent per month for the F160BW filter. Table 6.10 shows Table 6.10: Change in WFPC2 Throughput Over 30 Daysa. Filter PC1 WF2 WF3 WF4 ------------------------------------------------------------------------------ F16OBW -0.263 +- 0.030 -0.393 +- 0.051 F17OW -0.160 +- 0.011 -0.284 +- 0.005 -0.285 +- 0.006 -0.232 +- 0.006 F218W -0.138 +- 0.009 -0.255 +- 0.010 F255W -0.070 +- 0.007 -0.143 +- 0.009 F336W -0.016 +- 0.008 -0.057 +- 0.011 (-0.038 +- 0.018) (-0.043 +- 0.010) (-0.046 +- 0.008) (-0.047 +- 0.007) F439W -0.002 +- 0.007 -0.021 +- 0.010 (0.002 +- 0.014) (-0.022 0.007) (-0.023 0.009) (-0.023 0.007) F555W -0.014 +- 0.006 -0.016 0.008 (0.007 +- 0.013) (-0.007 0.007) (-0.009 0.009) (-0.008+- 0.008) F675W -0.001 +- 0.006 -0.001 0.006 (-O.020 +- 0.020) (0.001 +- 0.011) (0.002 +- 0.011) (0.004 +- 0.011) F814W 0.007 +- 0.007 0.003 +- 0.008 (0.013 +- 0.019) (-0.002 0.009) (-0.000 +- 0.009) (-0.002 +- 0.010) ------------------------------------------------------------------------------ a. Values in parentheses are from the omega Cen observations. the monthly decline in throughput based on this data. The values in parentheses are based on similar observations of the globular cluster omega Cen (NGC 5139; mean B-V ~ 0.7 mag). In general the values derived from the omega Cen data are in good agreement with the values derived from GRW+70D5824 data. A slight difference between the throughput declines for GRW+70D5824 and omega Cen might be expected due to differences in spectral shape, especially for filters like F336W which have a substantial red leak. However, even in the case of F336W the effect should be less than 0.01 mag based on SYNPHOT simulations. Figure 6.12 shows the throughput decline in all four chips as a function of days since the last decontamination for the F17OW filter. The contamination rate is remarkably constant during each decontamination cycle, and can be accurately modeled by a simple linear decline in throughput following the decontaminations, which appear to return the throughput to roughly the nominal value each month. While the contamination rates are similar for the three WF chips, the values for the PC are much lower. In addition to the monthly changes in throughput there is now evidence for a long-term variation in the F17OW data on the PC, where the throughput has increased at the rate of 4.8 percent +- 0.3 percent per year. This is evident in Figure 6.11, but is much clearer in Figure 6.12 where the data before and after Jan. 1, 1995 are compared, and in Figure 6.13 where the effect of the monthly decontamination is removed. The F160BW filter shows an even stronger trend but with larger uncertainties (i.e., an increase of 9.0 percent +- 1.7 percent per year). The WF chips do NOT show this effect, nor do the observations on the PC at longer wavelengths. One possible explanation of the throughput increase is that WFPC2 was flown with some initial contaminant on the PC1 optics which is slowly evaporating on-orbit. The pre-launch thermal vacuum test gave evidence of elevated contamination in PC1, which is consistent with this hypothesis. ISR 96-4 will describe detailed results of this monitoring (available from our site). Observers are advised to consult the STScI WFPC2 page for the latest information at the following address: http://www.stsci.edu/ftp/instrument_news/WFPC2/wfpc2_top.html Figure 6.11: Photometric Monitoring Data for WFPC2. Figure 6.12: Throughput for the Fl 7OW Filter Following Decontaminations. ------------------------------------------------------------------------------ CHAPTER 7: Observation Strategies In This Chapter... Observing Faint Targets Observing Bright Targets Observing Faint Targets Near Bright Objects Cosmic Rays Choosing Exposure Times Dither Strategies Pointing Accuracy CCD Position and Orientation on Sky Polarization Observations Observing with Linear Ramp Filters Emission Line Observations of Galaxy Nuclei 7.1 Observing Faint Targets For broad band filters the sky background will limit the detection of faint targets. For example, an 8 orbit observation in F555W gives a 5sigma detection limit at V=28.6 for an average sky level 23 mag arcsec^-2 in V. Note that the sky background is a strong function of position, especially for targets near the ecliptic; the sky level can vary from V=23.3 mag arcsec^-2 at the ecliptic pole to about V=20.9 mag arcsec^-2 on the ecliptic near the solar avoidance limit. (See Table 6.3 on page 113 for sky level as function of ecliptic coordinates.) If these higher sky levels would severely impact the science data, observers should consider specifying the special requirement LOW-SKY on the Phase II proposal. This parameter forces the observation to be made when the sky background is within 30 percent of the minimum value for the target. Note, however, that this will also reduce the number of HST calendar windows available to the observation, and so could result in scheduling delays or may even make the observation infeasible if there are other constraints such as ORIENTS. A minor decrease in the per-orbit visibility period also results from LOW-SKY, but for background limited programs this is a minor price to pay for the guarantee of a much lower background. In summary, LOW-SKY will reduce the sky background, but should only be used if the science goals require it. Note that LOW-SKY cannot be used for CVZ targets, as they imply mutually exclusive pointing constraints. Scattering of bright Earth light in the OTA can produce non-uniformities in the background which may hamper analysis of faint target images. Most often these take the form of diagonal bars of suppressed background light in several of the CCDS. These effects tend to occur for broad band filters when the OTA axis is about 25 degrees from the bright Earth. This effect is most often seen in observations of targets in the CVZ (continuous viewing zone), since the Earth limb is never very far from the OTA axis when observing in the CVZ. Figure 7.1 shows a typical case. LOW-SKY will eliminate this effect for non-CVZ targets, as it places the OTA axis more than 40 degrees from the bright Earth. Alternatively, one can place the target away from the CCD center to avoid these artifacts. Another option for reducing the sky brightness, is the special requirement SHADOW, which forces the observation to be made when HST is in the Earth's shadow. This usually has a large negative impact on the observing efficiency, and is recommended only when observing far-UV emission lines (e.g. Ly alpha and OI 1304A). Its primary goal is only to reduce geocoronal emission lines. Moreover, it does not attempt to minimize zodiacal emission, which dominates at visible wavelengths. 7.2 Observing Bright Targets Saturation is the primary concern when observing bright targets. The analog-to-digital converter will run out of bit codes at ~28,000 e- pixel^-1 in the ATD-GAIN=7 (default) setting, and at ~53,000 e- pixel^-1 in the ATD-GAIN=15 setting. Count levels above these are merely reported as 4095 DN in the data. Hence ATD-GAIN=15 is recommended for targets approaching 28,000 e- pixel^-1. The disadvantage of this setting is that the read noise is poorly sampled by this coarse digitization, and hence the read noise is slightly increased. At count levels above ~90,000 e- pixel^-1 charge will overflow the potential well of each pixel, and begin to bloom up and down the CCD columns. For example, this occurs in the F555W filter at about V=13.5 for a 10s exposure on the WFC, and at about V=13.0 on PC1. At very high count levels, above ~10^8 e- per CCD column, the charge bloom will reach the top and bottom of the CCD and flow into the serial registers. CLOCKS=YES will dispose of this charge as it reaches the ends of the CCD, and thus prevent it from leaking back into adjacent CCD columns. This exposure level corresponds roughly to a 10s exposure of a V=7 star in F555W. Note that CLOCKS=YES offers no benefit unless the bloom reaches the ends of the CCD, and that it may slightly compromise the bias and dark calibration. Moreover, CLOCKS=YES will result in anomalous exposure times; exposure times are rounded to the nearest integral second, minus a delay time of up to 0.25s for the Figure 7.1: Example of Scattered Earth Light. Scattered light contributes ~100 e- of background throughout this image. The camera spiders block some of this scattered light along CCD diagonals, hence forming "X" patterns and bars where the background is reduced by ~40 percent in this image. shutter to open. (See "Serial Clocks" on page 27 for further discussion of exposure time anomalies caused by CLOCKS=YES.) Besides setting ATD-GAIN=15, the PC CCD can be used to reduce saturation effects for stellar objects. The peak of the PSF will be spread over more pixels on the PC (vs. WFC), so stars can be exposed about 50 percent longer on the PC before saturation sets in. Note that the narrow band filters may be useful when observing very bright targets. For example, stars as bright as V ~4.4 can be observed without saturation in F502N using the PC at ATD-GAIN=15 with a 0.11s exposure time. 7.3 Observing Faint Targets Near Bright Objects The concerns here are similar to those for observing bright targets; saturation and blooming of the bright companion PSF must not impact the faint target. Also, one may need to consider subtracting the PSF of the bright object, and effects which limit the accuracy of that subtraction. If the, bright companion will saturate and bloom, it will be necessary to rotate the CCD so that blooming along the CCD columns does not obliterate the faint target. See Figure 7.8 on page 161 for an illustration of the bloom directions. It may also be useful to orient the field so that the OTA diffraction spikes from the bright companion (along diagonal lines on the CCDS) avoid the faint target. Table 7.1 summarizes ORIENTs which can be used to avoid CCD blooming tracks and OTA diffraction spikes caused by bright objects. For example, if a faint companion is at PA 60 degrees on the sky relative to a bright companion, it would be advantageous to observe on PC1 with ORIENT=PA + 45 degrees = 105 degrees. Ideally, some range in ORIENT would be specified to ease scheduling, hence "ORIENT=90D TO 120D" might be specified on the Phase II proposal. Note that "ORIENT=270D TO 300D" is also feasible, and should be indicated in the visit level comments. Table 7.1: ORIENTs for Avoiding Bloom Tracks and Diffraction Spikes. "PA" is the position angle of the faint target relative to the bright object. Note that ORIENT should be between 0D and 360D, so subtract 360 degrees, if necessary On proposal these are specified as, e.g., "ORIENT=231D TO 261D". CCD ORIENT ------------------------------------------------------------------------------ PC1 PA+30 degrees to PA+60 degrees, PA+210 degrees to PA+240 degrees WF2 PA+120 degrees to PA+150 degrees, PA+300 degrees to PA+330 degrees WF3 PA+30 degrees to PA+60 degrees, PA+210 degrees to PA+240 degrees WF4 PA+120 degrees to PA+150 degrees, PA+300 degrees to PA+330 degrees ------------------------------------------------------------------------------ If instead of observing a known companion, one is searching for companions, it is advisable to observe at several ORIENTs so that the CCD bloom track and OTA diffraction spikes will not hide possible companions. For example, three ORIENTS, each separated by 60 degrees, would give good data at all possible companion position angles. If PSF subtraction will be needed during data analysis, then the PC CCD may have some advantage, since it provides better sampling of fine undulations in the PSF. It may also be useful to obtain observations of a second bright star for PSF calibration, though these may be of limited utility since thermal effects and OTA "breathing" can modify the telescope focus, and hence the PSF, on timescales of less than one hour. Any such PSF star should be similar in color to the target, and should be observed at the same CCD position (within 1" and with the same filter. Sub-pixel dithering may also be useful, so as to improve sampling of the PSF (see "Dither Strategies" on page 155). Figure 7.2 illustrates the effect of OTA breathing, and periodic focus adjustments, on PSF subtraction. It shows the difference between an "in focus" PSF and one where the OTA secondary mirror has been moved by 54 micro meters. This amount of focus change is comparable to the range of OTA "breathing" effects (timescale < 1 hour), and the periodic (semi-annual) focus adjustments of the OTA. Each panel shows a different contrast setting; the percentages indicate the energy per pixel which is plotted as white, expressed as a fraction of the total (un-subtracted) PSF energy. For example, features which are just white in the "0.003 percent" panel contain 0.003 percent of the total PSF energy in each pixel. In other words, the feature labeled "a" is, in effect, ~10 magnitudes fainter than the PSF of the bright object, so that it may be very difficult to detect a "real" companion object ~10 magnitudes fainter than the bright object, at this distance from the bright object. In a real PSF subtraction situation, other effects including PSF sampling, noise, and pointing instability would further degrade the subtraction. (The elongated appearance of the residuals in the PSF core is due to astigmatism in PC1.) Table 7.2 gives some quantitative indication of the performance expected for PSF subtractions in the high signal-to-noise limit. It gives the magnitude of "star-like" artifacts remaining in the subtracted image, as a function of distance from the bright object, and magnitude m_bright for the bright object. The right-most column gives an effective magnitude limit imposed by artifacts from the PSF subtraction. These results are derived for the 5 micro meter focus shift described above, and are for PC1 and filter F555W. It may be possible to do somewhat better than these limits by subtracting accurate model PSFs, or by finding an observed PSF with matching focus. Table 7.2: Approx. PSF Subtraction Artifact Magnitudes and Magnitude Limits. Distance from Effective Magnitude of Effective Faint Object Bright Object Subtraction Artifacts Detection Limit (3sigma) ------------------------------------------------------------------------------ 0.1" m_bright+4.7 m_bright+3.5 0.3" m_bright+8.6 m_bright+7.4 1" m_bright+11.4 m_bright+10.2 3" m_bright+13.2 m_bright+12.0 ------------------------------------------------------------------------------ Large angle scattering may also impact identification of very faint objects near very bright ones. This scattering appears to occur primarily in the camera relay optics, or in the CCD. Hence, if a faint target is more than ~10" from a bright object (i.e. very highly saturated object), it would be advisable to place the bright object on a different CCD, so as to minimize large angle scattering in the camera containing the faint target. See section on "Large Angle Scattering" on page 99. Note also that highly saturated PSFs exist for PC1 in filters F439W, F555W, F675W, and F814W, and for F606W on WF3; these may be useful when attempting to subtract the large-angle scattered light. As of this writing TinyTIM does not accurately model the large angle scattering, and should be used with caution when analyzing highly saturated images (Krist 1996). Figure 7.2: Impact of OTA Focus Shift on PSF Subtraction. Each image shows the difference between an "in focus" and a 5 micron defocused PSF at different contrast settings. Numbers indicate the energy per pixel which is plotted as white, as a percentage of total energy in the un-subtracted PSF. Based on TinyTIM models for PC1 in F555W filter. It is generally unwise to place bright companions or other bright objects just outside the area imaged by the CCDS. The region of the focal plane just outside the CCDs (within about 6" of the CCDS) contains a number of surfaces which can reflect light back onto the CCDS, hence placing bright targets there can have undesired results. Also, the un-imaged "L" shaped region surrounding PC1 should be avoided, since incomplete baffling of the relay optics allows out-of-focus images of objects in this region to fall on the CCDS. Figure 7.3 illustrates various bright object avoidance regions near the WFPC2 field-of-view; the indicated avoidance magnitudes will produce 0.0016 e- s^-1 pixel^-1 in the stray light pattern for F555W. Figures 7.4 and 7.5 show examples of artifacts which can result from bright stars near the PC1 CCD. The report "A Field Guide to WFPC2 Image Anomalies" (WFPC2 ISR 95-06, available on the WFPC2 WWW pages and from help@stsci.edu) gives more detailed discussions of artifacts associated with bright objects, and their avoidance. Figure 7.3: Bright Object Avoidance Regions Near WFPC2 FOV. Figure 7.4: Example of PC1 "Direct" Stray Light Ghost. Figure 7.5: Example of PC1 "Diffraction" Stray Light Ghost. 7.4 Cosmic Rays Cosmic rays will obliterate ~20 pixels per second per CCD. It is imperative that two or more images be obtained at each pointing position, if these artifacts are to be removed from the data. The default action by the Phase II proposal processing software is to split exposures longer than 600s into two near equal parts, so as to allow removal of the cosmic ray tracks. The CR-SPLIT and CR-TOLERANCE optional parameters on the Phase II proposal allow observers to adjust this behavior. CR-SPLIT can be set to either DEF (default), NO, or a numeric value (0.0 to 1.0) giving the fraction of the total exposure allotted to the first sub-exposure of the pair. CR-TOLERANCE indicates the spread allowed in dividing the exposure, as a fraction of the total exposure time. For example, the default CR-TOLERANCE=0.2 allows the first sub-exposure to range from 0.3 to 0.7 of the total exposure. Setting CR-TOLERANCE=0 will force equal-length sub-exposures. The required degree of cosmic-ray avoidance will depend on the science goals of the proposal; observations of a single small target will usually suffer much less impact from cosmic rays than programs needing very "clean" data over a large area. Table 7.3 gives very rough recommendations for the number of sub-exposures for a given total exposure time. Note that splitting into many sub-exposures introduces additional overhead time and will increase the noise for "read noise" limited exposures (usually exposures in UV or narrow band filters), and hence one should not use more sub-exposures than are truly required by the science goals. Table 7.3: Recommended Exposure Splittings. ------------------------------------------------------------------------------ Total Exposure Rough Recommended Time (s) Number of Sub-exposures Programs with Sin- Wide-area Search gle Small Target Programs < 300 1 3 300 - 600 1 or 2 4 600 - 1600 2 or 3 4 1600 - 5000 3 5 5000 - 10000 4 6 > 10000 One exposure per orbit (2400s each) ------------------------------------------------------------------------------ 7.5 Choosing Exposure Times The choice of exposure time generally depends on the signal-to-noise ratio required to meet the science goals. This can be assessed using information in Chapter 6 or Appendix B herein, or by using the on-line WWW Exposure Time Calculator tool. However, when packing orbits, one must often compromise somewhat and decide which exposures to lengthen or shorten. Table 7.4 may be helpful in this regard. It shows the total time required to execute a single CR-SPLIT=NO exposure, excluding any time needed to change filters. Note that the most efficient exposure times are those whose length approaches or equals, but does not exceed, an integral number of minutes plus 40s. Figure 7.6 illustrates event timings during a typical 60s WFPC2 exposure, similarly, Figure 7.7 illustrates events during a (more efficient) 100s exposure. (See "Overhead Times" on page 29 for more information about exposure timings.) Figure 7.6: Event Timings During a 60s WFPC2 Exposure. All events, except shutter opening, start on 1 minute spacecraft clock pulses. Both the CCD clear and readout of each CCD require 13.6s. This 60s exposure, including the filter change, requires 4 minutes. Due to the various overheads, shortening or lengthening an exposure can have unexpected effects on the orbit packing. For example, shortening an exposure from 400s to 350s has no effect on orbit packing; they both require 9 minutes to execute (CLOCKS=NO, the default setting). On the other hand, shortening an exposure from 180s to 160s trims the execution time by 2 minutes (again CLOCKS=NO, the default setting). CLOCKS=YES may have some advantage in a long series of exposures whose lengths are 180s or somewhat greater. Each savings of 1 minute can add up to a few more exposures per orbit. The down side is that most calibrations are derived for exposures with CLOCKS=NO, so the calibration maybe slightly compromised. The largest calibration error is expected to occur on the dark current, where there may be a slight increase near the top and bottom of each CCD. In many situations this error may be acceptable, such as a small target near a CCD center, or broad band filter images where the sky completely dominates the dark current. CLOCKS=YES will have more impact on calibration of narrow filters, or situations requiring an extremely flat background. (Also, see "Serial Clocks" on page 27 for discussion of exposure time anomalies associated with CLOCKS=YES, though these are most important for exposures < 30s.) An exposure with CR-SPLIT=YES would simply require the total time for each sub-exposure as given by Table 7.4, again, plus any time needed to change filter before the first exposure. However, the default CR-SPLITting allows schedulers some latitude in dividing the exposures (CR-TOLERANCE=0.2 is the default) so the exact overheads are unpredictable. For example, a 700s exposure with CR-SPLIT=0.5 (the default) could be split into a pair of 350s exposures totaling 18 minutes, or a 300s and 400s exposure totaling 17 minutes. Table 7.4: Basic Time to Execute Single Non-CR-SPLIT Exposure. This includes time to prep the CCD, execute the exposure, and readout the CCDS. Times needed to change filter (1 minute), or insert a second filter (1 minute), are excluded. See "Overhead Times" on page 29 for more discussion and other overheads. Exposure Total Execution Time (min.) Time (s) CLOCKS=NO CLOCKS=YES (default) ------------------------------------------------------------------------------ 0.11 to 30 2 (not recommended) 35,40 2 2 50,60,70,80,100 3 3 120,140,160 4 4 180,200 6 5 230,260 7 6 300 8 7 350,400 9 8 500 11 10 600 13 12 700 14 13 800 16 15 900 18 17 1000 19 18 1100 21 20 1200 23 22 ------------------------------------------------------------------------------ 7.6 Dither Strategies There is no single observing strategy that is entirely satisfactory in all circumstances for WFPC2. One must consider cosmic rays, hot pixels (i.e. pixels with high, time variable dark count), spatial undersampling of the image, and large-scale irregularities such as the few arcsecond wide region where the CCDs adjoin. One strategy that can be used to minimize the effects of undersampling and to reduce the effects of hot pixels and imperfect flat fields is to dither, that is to offset the telescope by either integer or sub-pixel steps. The best choice for the number and size of the dithers depends on the amount of time available and the goals of the project. In the following we will address a few issues related to dithering: 1. Undersampling: Individual images taken with sub-pixel offsets can be combined to form an image with higher spatial resolution than that of the original images. A single dither from the original pixel position -- call it (0,0) -- to one offset by half a pixel in both x and y, (0.5,0.5) will produce a substantial gain in spatial information. On the other hand very little extra information is gained from obtaining more than four positions, if the standard four point dither is used, and if the telescope has successfully executed the dither. Therefore the recommended number of sub-pixel dither positions is between 2 and 4. 2. Hot Pixels: There are three ways to deal with hot pixels: correct using "dark frames" that bracket the observation, dither by an integer amount of pixels, or use a task such as "WARMPIX" within STSDAS to filter out the known hot pixels. Note that the integer dither strategy would ideally use six images, i.e. two CR-SPLIT images at each of three different dither positions. This is because in addition to hot pixels, low or "cold" pixels can be present and simple strategies selecting the minimum of two pixel values can fail. However, even four images (two each at two dither positions) will greatly aid in eliminating hot pixel artifacts. Cold pixels usually result from hot pixels in the dark calibration file which do not actually appear in the science data. 3. Cosmic Rays: Although dithering naturally provides many images of the same field it is better to take several images at each single pointing in order to remove cosmic rays. In principle, it should be possible to remove cosmic rays using only sub-pixel dithered data. At present, however, no publicly released software is available for this task. Hence we recommend obtaining two or more images (i.e. CR-SPLITing) at each position in the dithered sequence. For very long integrations it is convenient to split the exposure into more than two separate images. As an example, for two 2000s exposures, about 1000 pixels per chip will be hit in both images and will therefore be unrecoverable. Moreover, since CR events typically affect 7 pixels per event, these pixels will not be independently placed, but rather will frequently be adjacent to other unrecoverable pixels. 4. Accuracy of dithering: We do not yet have good statistics on the accuracy of HST dither offsets. During the Hubble Deep Field, nearly all dithers were placed to within 10 mas during +- 1.3" offsets and returns separated by multiple days), although in a few cases the dither was off by more than 25 mas, and on one occasion (out of 107 reacquisitions) the telescope locked on a secondary FGS peak causing the pointing to be off by approximately 1" as well as a field rotation of about 4 arcminutes. The software which was developed for the Hubble Deep Field is able to reconstruct images even for these non-optimal dithers, still gaining in resolution over non-dithered data. This software will be available in STSDAS in mid 1996 and is based on the variable pixel linear reconstruction technique developed by Fruchter and Hook (this procedure is also known as "dripping and drizzling"). The Richardson-Lucy-Hook non-linear deconvolution technique (STSDAS task ACOADD) can also use non-optimal dithers. The simplest way to schedule dithers is to use the options DITHER-TYPE=LINE (e.g. for two-point diagonal dithers) or DITHER-TYPE=BOX (for four-point dithers). Alternative approaches are to specify a spatial scan or to use POS TARGS. Note that when the WF3 is specified as an aperture, the POS TARG axes run exactly along the WF3 rows and columns. For the other chips, they only run approximately along the rows and columns due to the small amount of rotation between CCDS. For small dithers these rotations are not important. Some specific offsets allow one to shift by convenient amounts both the PC and the WFC chips. For instance an offset of 0.5" is equivalent to 5 WFC pixels and 11 PC pixels. The default DITHER-LINE spacing of 0.3535" along the diagonal is equivalent to shifts of (2.5,2.5) pixels for the WFC and (5.5,5.5) pixels for the PC. Note that large dithers will incur errors due to the camera geometric distortion which increases toward the CCD comers and alters the image scale by about 2 percent at the comers. Hence a 1.993" offset will be 20.3 WF pixels at the field center, but suffer a 0.4 pixel error at the CCD comers. Large dithers may also occasionally require a different set of guide stars for each pointing, thus greatly reducing the expected pointing accuracy (accuracy only ~1" due to guide star catalogue). One set of software presently available to handle the reduction of large dithers (which show the effects of geometric distortion) is the Fruchter-Hook code developed for reduction of the Hubble Deep Field images. For related articles on dither strategies, see the January, 1995 issue of the WFPC2 Space Telescope Analysis Newsletter and the February, 1995 issue of the ST-ECF Newsletter. For more information on dithering, please see "Questions about Dithering WFPC2 Observations" on the "Frequently Asked Questions" page of the WFPC2 WWW area. For detailed information on POS TARGS, please see the report "Dithering: Relationship Between POS TARGs and CCD Rows/Columns" obtainable from the WFPC2 WWW pages or help@stsci.edu. 7.7 Pointing Accuracy Some WFPC2 programs have critical target positioning constraints (i.e. the target must be as close as possible to a specified aperture). A sure way to meet such requirements is to include an interactive acquisition. However, INT ACQs are costly in terms of allotted orbits. The only alternative is to provide accurate target coordinates. This section provides some guidance in doing that, and indicates the level of residual error that can be expected. 7.7.1 Absolute Pointing Accuracy We have looked carefully at a sequence of images to assess the absolute pointing performance that HST delivers to WFPC2. The apertures used in the observations studied were either PC1, PC1-FIX, or WF2. The observed positions of stars on WFPC2 images were measured and compared with the proposed coordinates and apertures. Where necessary, coordinate and proper motion errors were accounted for (with the assumption that SAO catalog coordinates are exact they form the astrometric basis for the guide star coordinate system). The typical residual pointing error is 0.86", with 1.84" being the largest error seen. This study did bring out several easy-to-make target coordinate errors (which we corrected in the analysis, but which frequently dominated the pointing error), so we discuss these first. In a number of cases studied, the proposal coordinates were from the printed version of the Yale Bright Star Catalog. One problem is that the equinox 2000 positions in the BSC are given in the FK4 (Besselian) reference system. The proposal system assumes that equinox 2000 and later coordinates are in the FK5 (Julian) reference frame, and that earlier ones are in the FK4 frame. This can be overridden by specifying B2000 instead of J2000 for the equinox in the proposal. The latest digital version of the BSC (BSC5) is in J2000. The 1950 edition of the SAO catalog is in B 1950 (FK4), and a digital version is available for J2000 (FK5). An error of up to 1.5" can result from assuming BSC positions are J2000 instead of B2000 in the proposal. Another common problem with target coordinates is that they lack precision For example, in the BSC, RA is given to the nearest 0.1^s and DEC to the nearest arcsecond. This can cause an error of up to 0.75" in RA and 0.5" in DEC. The SAO coordinates have higher precision, 0.001s in RA and 0.01^s in DEC, and should be used when possible. A common error source is not specifying proper motion or specifying it in the wrong units. It is critical to follow the latest version of the proposal instructions on this because the units have changed. Even small proper motions are significant at the resolution of HST images. Residual pointing errors (after coordinate errors and aperture location changes) range from 0.26" to 1.84". The average is 0.93" and the median is 0.86". There are no obvious trends in any coordinate system. These are errors which cannot be accounted for by a proposer, being due to guide star position alignment uncertainties, and residual aperture location errors. Using Guide Star Catalog positions may help reduce the error between target and guide stars. Most of the targets used in this study were too bright to have true Guide Star Catalog positions. In summary, a target with good coordinates (and proper motion) referenced to the SAO catalog can typically be placed within 0.9" of a specified aperture. However, errors of around 1.5" occasionally happen. 7.7.2 Updates to Aperture / Coordinate Systems We note that on 11 April, 1994 an update was made to the spacecraft database which tells HST where to place targets relative to the FGSS. This impacted both the location of targets in the WFPC2 field-of-view, and the position reference frame in the image headers. The nominal (or intended) pixel locations of the apertures in the WFPC2 focal plane did not change. Only the (V2,V3) coordinates of the apertures changed, as their locations relative to the FGSs became better known. For example, PC1 and PC1-FIX are designated to be at pixel (420,424.5). Before April 1994, this aperture was thought to be at (V2,V3) = (4.95",-30.77"), which using the most current information was actually located at pixel (X,Y) = (459.8,377.3). Since April 1994, the aperture in the spacecraft's database has been at (V2,V3) = (1.87",-30.96) or, assuming the current best estimate is exactly correct, at (X,Y) = (414.9,428.1). Thus, for the same coordinates and aperture, the pixel position of a target in an image taken before April 1994 could be nearly 3" different from its position in later images, due to aperture updates. Similar corrections apply to all WFPC2 data taken before this date. This update also affects the position information placed in the image headers, which maps sky coordinates onto each individual CCD. Observations taken before April 11, 1994 have preliminary plate scales, rotations, and reference pixel locations in their image headers. Thus, the sky coordinates associated with a given pixel will be different for otherwise identical images taken before and after April 11, 1994, due to improvements in the aperture locations. The change is primarily an approximate 3" shift, as well as a small rotation. There is a 0.8 degree rotation for WF2, and smaller rotations for the other chips (0.28 degree in PC1, 0.46 degree in WF3, and 0.06 degree in WF4). The STSDAS task METRIC, which is used to assemble the 4 CCDs into a single image, uses this header information. Hence images taken before 11 April 1994 will not properly assemble into a single image. A new STSDAS task UCHCOORD will be available in the next STSDAS release (early 1996) to update the header group parameters to reflect the improved plate scale, shift, and rotations. This task will supersede the task UCHSCALE in the current STSDAS release, which can update only the scale, not the rotation or shift. We also note, that in April and May 1996 two updates were made to the (V2,V3) coordinate system. This update should be transparent to most observers. The purpose was to remove a slow drift in the position of WFPC2 in the HST focal plane; the largest change was 0.6". (See Table 3.15 on page 61 for details.) All the apertures are now thought to be correct to within 0.3", and updates in the near future should be small. 7.7.3 Pointing Repeatability The Hubble Deep Field (HDF) afforded an opportunity to study the repeatability of pointing over many images and acquisitions of the same field. The pointing appears to have been stable to better than 5 mas accuracy while taking many images of the same field without interruption over several orbits. The accuracy for full-up acquisition of the same field after slewing to other targets appeared to be ~10 mas typically, with occasional 20 mas errors seen. However, a few large errors were seen; in about 1 in 100 acquisitions the FGSs locked-up incorrectly resulting in a ~1" error. Other programs report similar 3 mas pointing accuracy if simple re-acquisitions are done between orbits. Approximately once per day a "full-up" acquisition is usually required (for engineering reasons) where the dominant FGS is fixed in position, but the sub-dominant FGS performs a spiral search for the guidestar and tracks wherever the star is found. On rare occasion these full-up acquisitions produce position errors of several hundred mas, and field rotations of up to ~0.1 degrees, relative to previous images of the same field. This may impact long sequences of exposures requiring half a day or more to execute. 7.7.4 Tracking Modes Three guiding modes are available: (1) Gyro Hold, (2) Coarse Track, and (3) Fine Lock. Fine Lock (PCS MODE FINE) is used by default, since use of Coarse Track may be harmful to the Fine Guidance Sensors. Use of Gyro Hold (PCS MODE GYRO) is not generally recommended, even for snapshot (SNAP) observations, since the pointing accuracy is only 14". Also the drift rate is 0.0014" s^-1 so exposures > 100s can result in smeared images. However, if the reduced pointing accuracy can be tolerated, and the exposures are only a few seconds or less, Gyro Hold can give a significant savings in the target acquisition overhead time. 7.8 CCD Position and Orientation on Sky During observation the target is placed at the aperture (PC1, WF2, WFALL, etc.) specified on the Phase II proposal. Locations of the principal apertures are shown in Figure 7.8 (Table 3.14 on page 61 gives a complete list of apertures; the V2,V3) system here is post 1996 day 127). The POS TARG special requirement can be used when a position offset is needed. The target is positioned with offset "POS TARG x,y", measured in arcseconds, from the specified aperture. The approximate directions (within 1 degree) of the POS TARG offsets are shown in Figure 7.8. The exact directions of the offsets are parallel to the rows and columns of the CCD on which the aperture is specified. There are small rotations (few tenths of a degree) between the CCDS. (For detailed information see "Dithering: Relationship Between POS TARGs and CCD Rows/Columns" obtainable from the WFPC2 WWW pages or help@stsci.edu.) It is often useful to explicitly specify the desired rotation of the WFPC2 field-of-view on the sky. This is specified in the Phase II proposal using the ORIENT special requirement. It is defined as the PA (measured from North through East) of the +U3 axis on the sky. Figure 7.8 shows the CCD orientation and aperture locations relative to the U3 axis. In effect, the sequence of events is to first move the target to the desired aperture, then offset by any specified POS TARG from the aperture, and finally to rotate the target "in place" on the CCDs to the desired ORIENT. Observers should try to specify all possible ORIENTs which would give the desired data, since having a range of values, or several ranges, will make the observation much easier to schedule. Often two ORIENTs separated by 180 degrees will both give useful data. Sometimes ORIENTs separated by 90 degrees will also give similar data. The ORIENT for any observation can be computed as follows: 1 . Obtain the Position Angle (PA) of the source axis on the sky, measured in the standard way, North through East. 2. Look at Figure 7.8 and decide what angle you want, measured clockwise, from the +U3 axis to the source axis. 3. Sum the angles in steps 1 and 2. 4. ORIENT must be between O and 360 degrees, so subtract 360 degrees, if necessary. The result is the ORIENT you should specify on the proposal. Another way to select the ORIENT, is to place Figure 7.8 on an image of the target, shift and rotate to get the desired alignment, and then simply measure the position angle of the +U3 axis relative to North. Note that the +V3 axis is quite different from the +U3 axis. They are exactly parallel, but oppositely directed. The +U3 axis is used for specifying orientation (ORIENT) in the proposal, while the +V3 axis is used in the data headers to indicate field orientation. Data header keyword PA_V3 gives the position angle of the +V3 axis on the sky. We now give two examples of how the POS TARG and ORIENT special requirements might be used. The first example (Figure 7.9) shows placement of a 100" long jet along the CCD diagonals in PC1 and WF3 (i.e. along the -U3 direction). The coordinates of the nucleus are given on the proposal. Aperture PC1 together with POS TARG +10, +10 are used to place the nucleus near the outer corner of PC1. We want to rotate the WFPC2 field-of-view about the nucleus so the jet is diagonal on PC1 and WF3. We can thus compute the desired orientation as ORIENT is defined as the Position Angle of the +U3 Axis on the Sky. Figure 7.8: ORIENT Definition, Aperture Positions, and CCD Alignments. "FIX" apertures are in same locations, unless otherwise indicated. Dashed lines show vignetted regions along CCD boundaries. Short lines and "X"s in outer CCD corners indicate directions of bloom and OTA diffraction spikes, respectively. Origin of the (V2, V3) system is at the origin of the plot axes, with V2 and V3 exactly along diagonal lines as marked. POS TARGs are offsets measured from the aperture specified on the proposal (PC1, WF2, WFALL, etc.); their directions are as indicated. CCDs have pixel (l,l) where the four CCDs overlap. ORIENT = (source PA on sky) + (desired source angle in field-of-view measured CW from +U3 axis) = 290 degrees + 180 degrees = 470 degrees - 360 degrees = 110 degrees On the Phase II proposal we would allow some range in the angle (to ease scheduling), hence "ORIENT 105D TO 115D" might be specified. The second example (Figure 7.10) shows placement of a galaxy across WF2 and WF3, with the nucleus on WF3 safely away from the vignetted region. Aperture WF3 together with POS TARG +20, 0 is used to place the nucleus near the outer edge of WF3. We want to rotate the WFPC2 field-of-view about the nucleus so the galaxy major axis is across WF2 and WF3. We can thus compute the desired orientation as ORIENT = (source PA on sky) + (desired source angle in field-of-view measured CW from +U3 axis) = 60 degrees + 315 degrees = 375 degrees - 360 degrees = 15 degrees On the Phase II proposal we would again allow some range in the angle (to ease scheduling), hence "ORIENT 5D TO 25D" would be specified. Note that "ORIENT 185D TO 205D" is also feasible, and should be indicated in the visit level comments. Note also, that WF3 and WF4 could be used with either "ORIENT 95D TO 115D" or "ORIENT 275D TO 295D" 7.9 Polarization Observations Polarization observations require three or more images with the polarizing filter spanning a large range of position angles on the sky. For WFPC2, this may be achieved by using different quads of the polarizing filter (each quad being oriented 45 degrees to the others), by rotating the spacecraft though different angles, or by a combination of these methods. Rotating the spacecraft through use of ORIENTs provides the simplest calibration, as a single polaroid can be used for all images. However, in practice, it will be the most difficult method to schedule. For further information see Biretta and Sparks (1995, WFPC2 ISR 95-01). We note that WFPC2 has significant instrumental polarizations which will make measurements on targets with less than 3 percent polarization difficult. The pick-off mirror introduces about 6 percent instrumental polarization. Furthermore, the pair of mirrors in the calibration channel, which is used to generate the polarizer flat fields, introduces ~12 percent polarization. In principle these effects can be calibrated out, but this has yet to be demonstrated. The polarizers are most effective in the range from 3000A to 650OA; this corresponds roughly to filters in the range F255W to F675W. At shorter wavelengths the transmission decreases sharply, and at longer wavelengths they cease to polarize the incoming light. Figure 7.9: Example of ORIENT and POS TARG Selection. (A) A jet at PA=290 degrees is observed using PC1 and WF3; the position of the nucleus is used for the target position. (B) The aperture is specified as "PC1" and the nucleus is placed near the outer corner of PC1 using "POS TARG +1O,+10." To place the jet across PC1 and WF3 "ORIENT 105D TO 115D" is specified. Figure 7.10: Example of ORIENT and POS TARG Selection. (A) A galaxy with major axis at PA=60 degrees is to be placed across WF2 and WF3. (B) The aperture is specified as "WF3" and the nucleus is placed near the outer edge of WF3 using "POS TARG +20,0." To place major axis across WF2 and WF3 "ORIENT 5D TO 25D" is specified. Note that "ORIENT 185D TO 205D" is also feasible. 7.10 Observing with Linear Ramp Filters The Linear Ramp Filters provide a narrow band ( deltalambda/lambda = 0.013 ) imaging capability which is continuously tunable from 3710A to 9762A. These are essentially a collection of narrow band interference filters whose central wavelength varies with position on the filter glass. The filter and aperture should be specified as LRF on the Phase II proposal, and the desired central wavelength should also be specified. The HST scheduling software will then select the target position so as to provide the desired wavelength. Note that it is not possible to choose between PC1 and WFC for the LRFS; one must use whatever CCD is automatically assigned by the scheduling software. If it is necessary to know which CCD will be used, observers can consult Table 3.7 on page 48 or Table 3.8 on page 49, or use the on-line LRF calculator tool on the WFPC2 pages. It is possible to use POS TARGs with LRF observations; the offsets are made from the default pointing for the specified wavelength. Observers should be mindful that the unvignetted field-of-view has a minimum size of ~10" in diameter, so that only small POS TARGs (< 4") should be used. While it is recommended that observers assume a 10" diameter field-of-view when using the LRFS, larger elongated (e.g. 15" x 10" targets can be sometimes be accommodated by placing the target major axis along the direction of the wavelength variation on the filter. This will result in a small reduction in throughput (i.e. small central wavelength offset) at the outer edges of the target. However, placing targets outside the central 10" of each ramp is strongly discouraged; outside the central 10" width the light will pass through more than one ramp segment, hence mixing light from different wavelengths, and making the data very difficult to calibrate. (See page 40 for further details on LRFS.) A common situation is one in which observers desire to make observations through an LRF filter, and then repeat the observation in a standard broad or narrow band filter at the same position on the CCD. The LRF Calculator Tool, available on the WFPC2 WWW pages, will tell observers the aperture (PC1-FIX, WF2-FIX, etc.) and POS TARG for any wavelength setting of the LRFS. Observers merely need to use this same aperture and POS TARG for the exposure through the other filter. If it is necessary to calculate the POS TARG manually, one can do this using the information in Table 3.7, Table 3.8, Table 3.14, and Figure 7.8. For example an LRF observation at 5034A would be made on WF2 at pixel (673.4, 235.7) (from interpolation by wavelength between X1 and X2, and between Y1 and Y2 in Table 3.7). These offsets are referred to the WF2-FIX aperture which is located (Table 3.14) at pixel (423.5,414). From Figure 7.8 we can deduce that pixel X direction is parallel to POS TARG "+Y" on WF2, and that pixel Y direction runs in the POS TARG "-X" direction. Using the pixel scale in section 5.10, "Optical Distortion", on page 103, we have POS TARG "X" = -0.09961 (235.7-414) = 17.76", and POS TARG "Y" = 0.09961 (673.4-423.5) = 24.89", hence POS TARG= +17.76,+24.89 would be requested for the non-LRF exposure. 7.11 Emission Line Observations of Galaxy Nuclei Saturation is a common problem for narrow band filter images of galaxy nuclei. Often the surface brightness of the emission line is estimated from ground-based images with 1" resolution; sometimes line fluxes are quoted for apertures several arcseconds in radius. However, at HST resolution, much of this flux may occur in a single unresolved spot at the galaxy nucleus, thus leading to saturated images. Observers should be cautious and consider this possibility when estimating exposure times. ------------------------------------------------------------------------------ CHAPTER 8: Calibration and Data Reduction In This Chapter... Calibration Observations and Reference Data Flat Fields Dark Frames Bias Frames Data Reduction and Data Products Pipeline Processing Fluxes and Standard Magnitudes Color Transformations of Primary Photometric Filters Cycle 5 Calibration Plan Cycle 6 Calibration Plan 8.1 Calibration Observations and Reference Data Standard calibration observations are obtained and maintained in the HST archive at the STScI, and can be obtained by external users using StarView. This includes those flat field, dark, and bias frames needed to operate the Post Observation Data Processing System (PODPS; sometimes called OPUS, and usually just called the "pipeline"), photometric calibration derived from standard star observations and the measured filter profiles, and derived determinations of the plate scale, distortion, and so on. The first set of these calibrations was provided to the STScI by the WFPC2 IDT from the Servicing Mission Observatory Verification (SMOV) and System Level Thermal Vacuum (SLTV) testing periods, and has been maintained and updated thereafter by the STScI with assistance from the IDT as part of the long term calibration program. For measurements requiring more precise calibrations, special calibration observations may need to be obtained as part of the observing proposal. Please consult the STScI WFPC2 Instrument Scientists for guidance if the routine calibration appears unlikely to support the requirements of a proposed observation. Individual GO programs requiring special calibrations must directly request these observations as part of their Phase I proposal. A database of laboratory characterizations of optical components, CCD sensors, filters, and the flat field channel has been collected to support the instrument calibration. On-orbit pointed calibrations require large HST resources, taking time that could otherwise be used for direct scientific observations. They can also be unsatisfactory due to the limitations of the available astronomical reference sources. For WFPC2, the inherent stability and uniformity of the CCD sensors, the well calibrated filters, the internal flat field calibration system, and an archive populated with flat field images obtained in SLTV prior to launch improve the scientific data analysis and productivity. Hence the need for on-orbit calibrations has been minimized. 8.2 Flat Fields The process of correcting for the effect of the variation in the sensitivity of the WFPC2 with field position is usually known as flat-fielding, or flattening. A "flat field" (an exposure of a spatially uniform source) needs to be observed through the telescope and desired filter. Real flat fields are always external; however, the WFPC2 has an internal calibration channel which produces a reasonably flat illumination pattern down to about 1800A. This channel is used to monitor and correct for changes in the flat fields. The instrument flat field is also coarsely monitored using the internal flat channel, using exposures called INTFLATS. Flat fields in narrow bandpass filters are obtained using the sunlit Earth (Target = EARTH-CALIB) as part of routine calibration in order to provide an absolute reference for the internal calibrations (and remove the low frequency effects of variations in the OTA illumination pattern). The Earth is an imperfect flat field target because it is too bright for the WFPC2 in the broad-band green and red filters. The rapid motion of the HST also creates streaks across the flat field images. The removal of the streaks requires the combination of multiple Earth observations with the streaks at different angles on the CCDS. An extensive discussion of the generation of Earth flat fields is available in Chapter 6 of the WF/PC-1 IDT OV/SV Report. The flat field calibration system works by imaging an illuminated diffuser plate onto the WFPC2 exit pupil (relay secondary) by means of an MgF2 lens. Two lamps provide optical and FUV illumination, providing a flat field which resembles the input beam from the OTA between 1600A and 10000A. During SLTV, flat fields were obtained using both the flat field calibration module and the WFPC2 optical stimulus (HST simulator), to generate a database of ratio images which link the internal flats to external flats. The external stimulus flats have been updated by comparison against on-orbit streak flats, obtained for a small subset of narrow band filters (F375N, F502N and F953N). These were obtained during the initial in-flight calibration of the instrument, to effectively calibrate low frequencies in the internal flat-fielding source. Eventually, sky flats for the broad photometric filters will be made by combining many frames. Flat fields have been obtained for all filters by combining information from SLTV test flats (which are good for all but the lowest spatial frequencies), and Earth flats in a limited number of filters (which fix the low frequency terms). The calibration channel monitors for time dependent changes in the flat fields, and none have been seen to date in the visible filters. Some redundancy is provided by the internal flat channel, but it does not form a part of the baseline calibration. During early 1996 flat fields for many filters were updated using an improved illumination pattern derived from large numbers of streak flats. Corrections were typically <= 1 percent, though the outermost comer of WF2 showed a 7 percent correction. The improved flats compared very favorably with sky flats. Note that the flat fields presently used in the pipeline are based on gain 14 data. The gain ratios are not constant from chip to chip, and therefore a small correction to photometric results derived from gain 7 data should be applied (see Table 4.3 on page 82). (See Biretta 1995 for further discussion of WFPC2 flat fields; also see the HST Data Handbook.) 8.3 Dark Frames Dark frames are long exposures that are taken with no light incident on the CCDS. They are used to detect CCD counts (the dark current) caused by thermal generation (at the interfaces between the silicon and oxide layers), and the rate of charged particle and secondary radiation events. Estimated dark current and cosmic ray event rates are given in section 4.8, "Dark Backgrounds", on page 71 and section 4.9, "Cosmic Rays", on page 75, respectively. Observers are cautioned that the calibration provided by the pipeline may not use the most up-to-date dark frames (partly because constructing these may involve data taken after the observation). This matters because new hot pixels are being generated all the time. Data to construct improved super-darks is available in the archive with at least 5 long dark exposures being taken weekly. 8.4 Bias Frames Bias frames are readouts of the CCDs without an exposure (so the dark current is negligible). They are used to measure the DC offset built into the signal chain which ensures that the ADC input is above zero. They fix the constant that is subtracted from the raw data prior to dark subtraction and flat fielding in Routine Science Data Processing (RSDP). 8.5 Data Reduction and Data Products The routine processing of WFPC2 science data consists of the pipeline functions described below. The resulting images will be available on magnetic tape in FITS format, and as greyscale printouts. The reformatted raw data will also be available, along with the relevant calibration data. The STSDAS Calibration Guide should be consulted for a more complete description than the summary presented here. * The following data are supplied to observers on FITS tapes: * Edited Image and DQF (uncalibrated): doh, .qoh * Standard Header Packet: .shh * Extracted Engineering Data and DQF: .x0h, .q1h * Trailer File (ASCII file): trl * Calibrated Image and DQF: .c0h, .c1h In addition, a histogram file used for monitoring of the signal chain (.c2h file), and a calibration file which gives the throughput curve used in populating the photometric keywords (.c3t file) are included. Further data reduction and analysis can be performed under the STScI's science data analysis software system (STSDAS). Standard routines are available, operating under IRAF, for the analysis of data for image photometry, spectral analysis, astrometry, and the generation of the calibration data files. 8.6 Pipeline Processing The pipeline processing of WFPC2 data sets reformats the telemetry received from HST into group FITS format images, generates headers containing a large number of keywords taken from both the HST and WFPC2 telemetry streams, in addition to various STScI ground system databases, and applies the calibration described below. This calibration is done with a module known as "CALWP2" which is written in the IRAF SPP language and is available, in identical form, to users of the STSDAS system. Therefore, off-line recalibration of observations is fairly easy, and will use the same program as the PODPS system. Documentation is available in the HST Data Handbook, and the STSDAS User's Guide. CALWP2 performs the following operations if required by the observation: * ADC correction * Bias level removal * Bias image subtraction (depending on the gain channel in use) * Dark image scaling and subtraction * Shutter shading correction * Flat field image correction * Population of various photometric calibration keywords In addition, the following conditions are flagged in the Data Quality File (DQF): * Transmission failures and other possible failures * Known bad pixels (e.g. blocked columns) * Pixels at or above the maximum A/D converter level (i.e. saturated) * Bad pixels in reference images 8.7 Fluxes and Standard Magnitudes The pipeline calibrated data are not flux calibrated and the data are in units of Data Numbers (DN). However a flux calibration is implicit in the header. To get flux density, multiply DN by the value of the keyword PHOTFLAM in the calibrated (.c0h) science header file, and divide by the value of the keyword EXPTIME. The magnitude of an object can be determined using the photometric zero-point keyword PHOTZPT as . m = -2.5log_10(PHOTFLAM x (DN/EXPTIME)) + PHOTZPT where m is in the STMAG system which is based on a spectrum with constant flux per unit wavelength set to roughly match the Johnson system at V. The more conventional systems are based on Vega's spectrum. Table 8.1 was generated using SYNPHOT to provide rough conversions to the Johnson UBVRI and Cousins RI systems. Typical uncertainties are 5 percent, and probably much worse for the U filter. The correction depends on the spectrum of the object, hence the table was generated using a wide range of Bruzual models. For example, to convert to the Cousins I band for an object on WF4, get PHOTZPT=-21.1 and PHOTFLAM=2.6044 x 10^-18 from the header. Then convert from WFPC2 counts to magnitudes in Cousins I using: I_c = -2.5log_10(2.6044 x 10^-18 x (DN/EXPTIME)) - 21.2 - 1.21 Note that the Cousins I filter is much closer to the F814W filter than Johnson 1, as shown by the nearly constant correction as a function of spectral type (i.e. color term). Table 8.1: Conversion from STMAG to Johnson UBVRI and Cousins RI. U-F336W B-F439W V-FS55W R_J-F675W I_J-F814W R_C-F675W I_C-F814W ------------------------------------------------------------------------------ O5V 0.53 0.67 0.05 -0.67 -1.11 -0.71 -1.22 B0V 0.46 0.66 0.05 -0.67 -1.13 -0.70 -1.22 A0V -0.08 0.67 0.02 -0.68 -1.22 -0.67 -1.21 F2V -0.03 0.62 -0.00 -0.69 -1.28 -0.63 -1.22 G0V -0.02 0.58 -0.01 -0.70 -1.31 -0.60 -1.23 K0V -0.10 0.53 -0.01 -0.69 -1.32 -0.58 -1.23 M0V -0.04 0.43 -0.00 -0.78 -1.48 -0.54 -1.22 M6V 0.05 0.29 -0.03 -1.05 -1.67 -0.56 -1.21 ------------------------------------------------------------------------------ This procedure will provide typical accuracies of about 0.05 mag (worse in the UV). More accurate photometry will require a variety of corrections (e.g., CTE effect, contamination and red leaks for the UV filters, variable gains on different chips, color terms, geometric distortions) which are discussed in detail in Holtzman et al. (P.A.S.P., 1995b). 8.8 Color Transformations of Primary Photometric Filters The WFPC2 UBVRI system is fairly close as regards effective wavelengths to the Johnson UBVRI system, but cross-calibration is necessary to convert to Johnson magnitudes. See the IDT OV/SV Report and Harris et al., A.J. 101, 677 (1991) for examples in the case of WF/PC-1. Figures 8.1 through 8.5 show the results of regression fits between these two systems on the main sequence stars in the Bruzual, Persson, Gunn, Stryker atlas that is installed in the calibration database system (CDBS). These fits should be used with caution for quantitative work. The zero-points in all cases are defined so that Vega's spectrum integrated over the bandpass is exactly magnitude zero (VEGAMAG in XCAL). The zero-points of the canonical Johnson-Cousins system differ from this by up to 0.02 magnitudes. The zero-points thus defined for the HST filters do not coincide with the STMAG definition used in the previous section. In addition, the ground based filter curves used, which are taken from Bessel (P.A.S.P. 102, 1181), give a good approximation to the standard Johnson-Cousins system, but are not as accurate as taking Landolt's curves and applying his color corrections to transform to the standard system. The latter procedure was used to derive the transformations Figure 8.1: F336W-F439W against Johnson U-B for the BPGS atlas of MS dwarf spectra. The change in slope in the transformation for U-B greater than about 0.1 is due to red leak in the F336W filter. For hotter stars, the transformation is quite linear. Figure 8.2: F439W-F555W against Johnson B-V. The residuals from the best linear fit are quite similar to those that apply if F569W (instead of the preferred F555W) is chosen for a WFPC2 passband. Figure 8.3: F555W-F814W against Johnson V-I_C. The residuals from the best linear fit are generally very small. This particular color combination is widely used. Figure 8.4: F555W-F675W against Johnson V-R_C. The residuals from the best linear fit are somewhat larger for blue stars than those that apply if F569W is chosen. given in Holtzman et al. (P.A.S.P. 1995b), which also discusses the changes in the transformations that result from source spectrum variations (such as metallicity and gravity effects). Figure 8.5: F675W-F814W against Cousins RC-I_C. The residuals from the best linear fit are similar to those that apply if F791W is chosen for a WFPC2 I passband. For spectral type M8V and later (not shown) the transformation will not work as well. 8.9 Cycle 5 Calibration Plan A summary of the Cycle 5 calibration plan follows as a general guide to the calibration and monitoring program in place for WFPC2. The full proposals are available through STScI's proposal status web page. http://presto.stsci.edu/public/propinfo.html The data that the calibration and monitoring program produces has no proprietary period and is available through the HST archives In each case the proposal ID is given followed by its title, purpose, description, accuracy goals and products produced. Calibration information obtained to date consists primarily of the System Level Thermal Vacuum (SLTV) tests, the initial on-orbit tests conducted in SMOV, and the Cycle 4 calibration. These tests have shown that the instrument is stable with some important exceptions and have provided an initial calibration sufficient for routine processing of most data. The Cycle 5 calibration is designed to enable users to maximize the scientific usefulness of their data, while at the same time minimizing the use of spacecraft time. This is done by designing efficient proposals that: * A) Improve the existing calibration - in particular towards the goal of 1 percent absolute photometric accuracy. * B) Assess the accuracy of the existing and new calibrations. * C) Recalibrate important known time variable features of the instrument. * D) Calibrate some important instrumental effects that are not well understood * E) Monitor the instrument and telescope to ensure that no new problems or variability in their performance are missed. * F) Maintain the instrument in a healthy state and ensure that in the event of partial instrumental failures, the calibration can be maintained when possible. The calibration of the instrument must be seen in a larger context than simply preparing reference files for a pipeline reduction and assessing the errors in them. Several calibrations (such as geometric distortion, CTE correction, PSF calibration, chip-to-chip alignments, polarization calibration) are very important to some observers. Yet they are not included in the pipeline. Other things frequently need to be done to the data after it is ADC, bias, dark, and flat field corrected, with a photometric calibration included in the header. These other calibrations are made available to users through this Instrument Handbook, journal publications, instrument science reports, and postings linked to the Institute's WFPC2 home page. The address is: http://www.stsci.edu/ftp/instrument_news/WFPC2/wfpc2_top.html A list of the most important calibrations consists of the following items: 1. Photometric zero-point. This involves converting DN values to flux units. 2. Photometric transformations. This involves converting DN values to magnitudes in standard systems. Two separate photometric calibrations can be used for this, a direct approach and a synthetic approach. 3. Photometric temporal variations. This is particularly important in the UV where significant variability is seen. 4. Photometric spatial variation (flat fields and CTE). 5. Dark current. This includes its time variability (hot pixels). 6. Bias. 7. Analog-to-Digital converter errors. 8. PSF. This is crucial for PSF fitting photometry, PSF subtraction, PSF modeling, and deconvolution efforts. Because PSF subtraction of very saturated sources is specialized to a few very diverse programs, PSF calibration in the image halo (beyond about 0.5 arcsecond) is not supported and must be requested with the program as a special calibration. 9. Polarization calibration. 10. Geometric calibration. The program consists of 15 proposals which use a total of 63 orbits of spacecraft time (to be compared to a total of about 1550 orbits of approved GO time). The proposal summaries and their associated RPS2 files largely speak for themselves. Table 8.2 lists all of the proposal numbers, titles, the schedule for the calibration execution, an indication of whether the output forms part of the pipeline data reduction (CDBS) or provides other information, usually documented in Instrument Science Reports (ISR), the approximate calibration accuracy expected (see the summary forms for the interpretation of these numbers, because they are almost meaningless without a context), and the primary areas from the above 10 calibration types they address and in what ways (A-F from the above list). Therefore, if you are interested in a particular calibration, look for its number in the last column of the table, and then refer to the subsequent proposal summaries for more details. Table 8.2: Summary of Cycle 5 Calibration Plans. External ID Proposal Title Schedule Results Accuracy Time Notes (percent) (orbits) ------------------------------------------------------------------------------ 6179 Photomet. Zero. Late 95 CDBS 1 1ABE, 2AB 6182 Photomet. Trans. 9/95, 3/96 CDBS 2 6 2ABE 6183 Decontamination 1x per 4 wks. ISR N/A 0 F 6184 Photometric Mon. 2x per 4 wks. ISR 1 24 3E 6186 UV Throughput Early in Cyc.5 CDBS 10 6 1AB, 3C 6187 Earth Flats Continuous CDBS 1 0 4ABE 6188 Darks Weekly CDBS 6 0 5ABC 6189 Visflat Monitor 2x per 4 weeks ISR 0.6 0 4E 6190 Internal Flats Early Cyc. 5 CDBS 0.6 0 4F,7E 6191 UV flats 2x in Cyc. 5 ISR 2 0 4ABE 6192 CTE Calibration Early Cyc. 5 TIPS < 1 4 4ABD 6193 PSF CTE+2m TIM 10 5 8ABD 6194 Polarizers+Ramps TBD CDBS 3 + 2 8 9DE, 1AB 6195 Flat field Check Late 95 CDBS 1 2 4B 6250 Internal Monitor 2x per week ISR N/A 0 5,6, 10F ------------------------------------------------------------------------------ TOTALS 63 a. Under Notes column, letters and numbers are keyed to lists in text. 6179: Photometric Zero-point * Purpose: Set synthetic zero points of all WFPC2 filters. * Description: GRW+7OD5824 is observed through all filters not in the photometric calibration monitors and longward of F336W (inclusive). It is observed in both PC1 and WF3. This data-set is directly comparable to the corresponding results for Cycle 4 (program 5572). Because of possible errors in the spectrophotometry for this target, and in order to check synthetic color transformations over a reasonably wide range of colors, the observations are repeated with 3 red standards, and 2 other blue standards in WF3 only. This time, most of the 18 broad and medium bandpass filters longward of F336W are included (restricted to 1 orbit/target). If CTE calibration fails, this proposal may need to be run with preflash (and take more time or do less targets). This proposal is needed by all GO proposals that want to do quantitative photometry at the few percent level. * Accuracy: Overall discrepancies between the synthetic photometric products and the results of this test should be reduced to 1 percent rms. Part of the point of the test is to measure this accuracy, which will largely depend on the accuracy of the spectrophotometric calibration sources. * Products: After pipeline processing, each image will be reduced by aperture photometry to measure the throughput of each filter. These numbers can be directly compared to SYNPHOT predictions. Systematic differences will be corrected in the throughput database by tweaking the filter normalizations (already done for the primary target in Cycle 4), overall system response (which is quite uncertain particularly longward of 8000A), and finally bandpass shapes. Residual differences will give an idea of the intrinsic accuracy of the calibration. 6182: Photometric Transformation * Purpose: Update photometric transformations to Johnson-Cousins system. * Description: A photometric standard star field in omega Cen is observed twice once at the September 1994 orientation and once rotated by 180 degrees (to correct to first order for residual CTE effects). All broad and medium bandpass filters are used. Based on Cycle 4 program 5663, this proposal also gives a check on the long term full field photometric stability of the instrument. * Accuracy: Independent of the synthetic photometry results this test gives direct transformations to the Johnson-Cousins system for a wide range of source colors (but excluding very blue stars). These transformations should be accurate to 2 percent. The stability of these transformations will be measured to the sub-percent level (because then errors in the ground based photometry do not enter significantly). * Products: A comparison with the corresponding Cycle 4 monitor (which ran monthly) will verify the photometric stability of the camera. Direct transformations to the Johnson-Cousins photometry system can be derived for all filters. The observations also provide a check of the astrometric and PSF stability of the instrument over its full field of view. 6183: Decontamination * Purpose: Remove UV blocking contaminants from CCD windows. * Description: A sequence of observations is defined that is run twice - once before and once after a DECON when the CCDs are heated to +20 degrees C for 6 hours. The sequence consists of 2 bias frames at each gain setting, 5 1800 second darks, 2 INTFLATS through F555W at each gain, and two K-spot images. The observations are arranged so that the first sequence occurs about 33 hours before the DECON, and the second follows it by 12 hours. The proposal is run every 4 weeks. This is based on Cycle 4 program 5568 with bias frames added. * Accuracy: This proposal is mainly designed to test aliveness, and monitor the instrument to ensure that no untoward effects from the DECON have occurred. It will identify all hot pixels that are annealed by the DECON. * Products: The data is examined and checked for anomalies. The dark frames are processed to yield plots showing the growth and annealing of hot pixels. 6184: Photometric Monitor * Purpose: Monthly external check of instrumental stability. * Description: GRW+7OD5824 is observed through F17OW in all four chips. It is then observed in one chip for filters F160BW, F185W, F218W, F255W, F300W, F336W, F439W, F555W, F675W, F814W to fill out the orbit. The chip chosen is changed each month, so each chip is used three times during the year. One extra F555W exposure is taken through the PC to allow for focus monitoring. The proposal is run once before and once after each monthly decontamination, with the same chip selected. Based on Cycle 4 programs 5629 + 6143 + (5563) + 5564. * Accuracy: Overall discrepancies between the results of this test should be less than 1 percent rms. The point of the test is to measure this variation. * Products: After pipeline processing, each image will be reduced by aperture photometry to measure the throughput of the filters. These numbers can be directly compared to results for previous months. This will allow the long term performance of the instrument to be checked for changes, and verify that the decontaminations are satisfactory. This proposal will be run every 4 weeks in association with the DECON. Results reported in Baggett et al., WFPC2 ISR 96-02 and Whitmore, et al., WFPC2 ISR 96-04 6186: UV Throughput * Purpose: Update SYNPHOT database for UV throughput. * Description: GRW+70D5824 is observed shortly before and after a DECON through all of the UV filters in each chip and through F16OBW crossed with F130LP, F185LP and F165LP (where applicable) to determine the wavelength dependence of the throughput across the bandpass (hence color terms). Based on no particular Cycle 4 program, this program is designed to characterize better the spectral response curve in the UV, and the spectral shape introduced by the contamination. * Accuracy: Overall discrepancies between the updated synthetic photometric products and the results of this test should be 1-2 percent rms. This does not mean that the UV throughput will be known to this accuracy, primarily because of uncertainties in the flux calibration of the standard used (5 percent), uncertainties in the UV flat fields (maybe 3 percent near the chip center), and time dependent contamination corrections (3 percent error), and uncertainties in the CTE correction (2 percent). The derived UV absolute photometric accuracy at the center of the chips should therefore be about 10 percent. * Products: After pipeline processing, each image will be reduced by aperture photometry. The throughput curves and their normalizations can be updated by trial and error. 6187: Earth Flats * Purpose: Generate flat fields. * Description: 4 sets of 200 Earth streak-flats are taken to construct high quality narrow-band flat fields with the filters F375N, F502N, F656N, and F953N. Of these 200 perhaps 50 will be at a suitable exposure level for de-streaking. The resulting Earth superflats map the OTA illumination pattern, and will be combined with SLTV data (and calibration channel data in case of variation) for the WFPC2 filter set, to generate a set of superflats capable of removing both the OTA illumination and pixel-to-pixel variations in the flat field. In addition, more limited observations are made in F16OBW and the broad bandpass photometric filters. The Cycle 4 plan is being largely repeated except: (1) UV filters are dropped because the measurement is generally only of the red leak. (2) F16OBW is retained in order to check for developing pinholes. (3) Crossed filters used as neutral densities are eliminated (illumination pattern is wrong). (4) An attempt will be made to schedule some broad bandpass measurements on the dark Earth. This is based on Cycle 4 programs 5570 + 5571 + 6142. * Accuracy: Overall accuracy of the flats derived from this test and the corresponding Cycle 4 observations should be below 1 percent RMS. Discrepancies between the results of this test and those from Cycle 4 should be 1 percent RMS. Differences between the two datasets analyzed separately will measure the flat field variability to this level. This data, together with the flat field check proposal should enable similar accuracy in the broad bandpass flats. * Products: This proposal provides medium and narrow bandpass streak flats which can be combined with the high frequency information in the TV flats to yield accurate flat fields. The ratio of the TV and derived flats provides a correction that can also be applied to the TV broad bandpass data. We may also get broad band flat fields directly for comparison from the sky or from this proposal's exposures on the Earth's shadow. 6188: Darks * Purpose: Provide dark frames for pipeline reduction, and hot pixel lists. * Description: 5 dark frames are taken every week to provide ongoing calibration of the CCD dark current rate and to monitor the characteristics and the evolution of hot pixels. Over an extended period, these data provide a monitor of radiation damage to the CCDS. The dark frames will be obtained with the CLOCKS=OFF. In addition, 4 darks are taken per month with CLOCKS=ON (although there is no effect such as amplifier glow expected from running the serial register for these CCDS). Changes from the Cycle 4 program 5562 are that each group of 5 darks is constrained to be taken within a 2-day period, in order to simplify the data analysis (because fewer new hot pixels are involved), and the CLOCKS=ON exposures have been added. * Accuracy: The accuracy of the super-dark computed from these data depends on the number of frames combined. The present practice is to combine them in groups of 10 frames for pipeline super-darks. This gives a median signal-to-noise of 16 and higher signal-to-noise on hotter pixels than the median (with the somewhat shaky assumption that the dark noise is Poisson). This means that the residual systematic error after super-dark subtraction on an 1800-second exposure is about 3 electrons - much less than the read noise. In principle, this residual can be further reduced to 0.4 electrons if a super-dark is generated from all of the dark frames, with suitable masking based on hot pixel lists. * Products: The data is grouped into sets of 10 frames every two weeks. These are combined into super-darks for use in the pipeline. In addition, hot pixel lists can be generated with a time resolution of one week. 6189: VISFLAT Monitor * Purpose: Monitor internal flat fields of instrument. * Description: All use of the VISFLAT channel is concentrated in this proposal. It is based on Cycle 4 programs 5568 and 5655. The program takes one complete set of exposures using the visible calibration channel lamp (VISFLATS) at the start of the cycle through all visible filters. Monthly, VISFLATS will be obtained with the photometric filter set (F336W, F439W, F555W, F675W, and F814W) both before and after the DECON. A monthly VISFLAT exposure with the Wood's filter, F16OBW, allows its visible transmission to be monitored. Two monthly VISFLAT exposures obtained through the LRF (FR533N), one at each gain, provide a monitor of the ADC's performance. The VISFLAT exposures should be packed together to optimize use of each lamp-on cycle. * Accuracy: The internal flats, when well exposed, are each accurate to 0.6 percent in terms of the pixel-to-pixel (high frequency) variations in the CCDS. Thus, high frequency flat field stability can be verified to 1 percent. When the results from several filters are combined it will be possible to check that the CCDs are indeed relatively stable to much better than 1 percent. * Products: The complete filter-set sweep will be compared to the corresponding Cycle 4 data-set primarily to verify that none of the filters are developing problems. Ratios of these flats will primarily indicate the stability of the channel itself, unless there are strong variations from filter to filter. Unless time dependence in the filters is seen, it is likely that flat fields for pipeline calibration will continue to be made by combining Earth-flat, SLTV-flats, and eventually sky-flats. The bimonthly photometric filter-set observations will be used to monitor WFPC2's flat field response and to build a high S/N flat field database (primarily useful in tracking any changes in the pixel-to- pixel response of the instrument and any possible long term contamination induced changes). Histograms generated from the ramp filter flats will be used to trace the ADC transfer curve. F16OBW can be checked monthly for pinholes. Results reported in Stiavelli and Baggett, WFPC2 ISR 96-01. 6190: Internal Flats * Purpose: Provide backup database of INTFLATS, in case VISFLAT channel fails. * Description: Based on Cycle 4 program 5568. A complete set of illuminated shutter blade flats (INTFLATS) is taken close to the complete set of VISFLATS. Each filter is exposed on each shutter blade (A or B) at each gain setting (7 or 15). Thus, there are 4 exposures per filter which should be uninterruptedly sequenced as (A7, A15, B15, B7). There is a possible concern on thermal control, where an out-of-limit condition was almost reached when 5568 was run. This will be avoided by spacing the exposures suitably. * Accuracy: The signal-to-noise per pixel is similar to that obtained in the VISFLAT program (0.6 percent) but there are much larger spatial and wavelength variations in the illumination pattern. As a result, this data-set will not form any part of the pipeline calibration. This baseline is necessary in case the VISFLAT channel fails, and there are temporal variations in the camera flat fields at the 1 percent level. The test does give a good measurement of the gain ratios and their stability, which should be accurate to much better than 1 percent, when all of the data is analyzed. * Products: INTFLATNISFLAT ratios can be generated if there is a failure in the calibration channel. Gain ratios and stability will be assessed. Results reported in Stiavelli and Baggett, WFPC2 ISR 96-01. 6191: UV Flats * Purpose: Use UV calibration channel to monitor long term internal UV stability. * Description: UV flat fields will be obtained with the calibration channel's ultraviolet lamp (UVFLAT) using the limited FUV filter set (F122M, F170W, F160BW, F185W, and F336W) twice in the cycle immediately after a DECON. The UV lamp is degrading with use, so its use must be minimized. The UV flats will be used to monitor the FUV flat field stability and the stability of the Wood's filter, F160BW, by using F170W as a reference. The VISFLAT/UVFLAT ratio from the F336W filter will provide a diagnostic of the UV flat field stability, and tie the UVFLAT and VISFLAT flat field patterns together. This program represents the entire use of the UV lamp in Cycle 5. This proposal is based on the Cycle 4 program 5568, but with two extra filters (F122M and F185W). * Accuracy: Should verify stability of the UV filters and flat field to 2 percent. The overall flatfield response is not measured because the lamp output is not uniform, and temporal variations in throughput are not measured because the lamp output varies. * Products: Ratio images with the corresponding data-set from Cycle 4 and future cycles will verify the UV flat field stability. 6192: CTE Calibration * Purpose: Calibrate CTE effect for a range of star brightness and background. * Description: The crowded omega Cen field is observed for 40 seconds through F555W with gain 7 and preflashes of 0, 5, 10, 20, 40, 80 and 160 electrons. As a gain check and calibration, it is observed at the same orientation with gain 15 twice at preflash levels of 0 and 160 electrons. The whole sequence is repeated with filter F814W. Then a whole orbit is filled with 1 second exposures, in order to investigate the effect of CTE on low signal level stars (but with high accumulated signal-to-noise). This last orbit is repeated with a preflash of 40 electrons. This is based on Cycle 4 programs 5645, + 5646 + 5659. * Accuracy: As this test is a differential measurement of the CTE slope, it should be very accurate (much better than 1 percent). As a large number of stars are involved, and the photon noise on each measurement is of order 1 percent, the slope derived should be much more accurate. The largest remaining uncertainty will be the absolute level of the slope, not differential effects caused by varying background. * Products: After pipeline processing, each image will be reduced by aperture photometry to measure the stellar brightness and how it depends on the preflash level. This is a differential measurement and gives no direct information about the slope of the CTE effect at high background levels. The absolute CTE can be estimated from 5646 (already run once with a raster in this field). 6193: PSF Characterization * Purpose: Provide a sub-sampled PSF over the full field to allow PSF fitting photometry. * Description: This proposal measures the PSF over the full field in photometric filters in order to update the TIM and TinyTIM models and to allow accurate empirical PSFs to be derived for PSF fitting photometry. With one orbit per photometric filter, a spatial scan is performed over a 4x4 grid on the CCD. The step size is 0.025 arcseconds. This gives a critically sampled PSF over most of the visible range. The crowded omega Cen field is used. 40 sec images are taken through each of the photometric filters (F336W, F439W, F555W, F675W, F814W). Data volume will be a problem, so special tape recorder management will be called for. This is based partially on Cycle 4 program 5575, which used the same field. The proposal also allows a check for sub-pixel phase effects on the integrated photometry. * Accuracy: The chosen field will have hundreds of well exposed stars in each chip. Each star will be measured 16 times per filter at different pixel phases. In principle, the proposal therefore provides a high signal-to-noise critically sampled PSR. This would leave PSF fitting photometrists in a much better position than now where pixel undersampling clearly limits the results. The result will be largely limited by breathing variations in focus. It is hard to judge the PSF accuracy that will result. If breathing is less than 5 microns peak-to-peak, the resulting PSFs should be good to about 10 percent in each pixel. PSF fitting results using this calibration would, of course, be much more accurate. In addition, the test gives a direct measurement of sub-pixel phase effects on photometry, which should be measured to much better than 1 percent. * Products: This program provides sub-sampled PSFs for photometry codes, data for comparison with PSF codes, and measurement of pixel phase effect on photometry (sub-pixel QE variations exist). 6194: Polarization and Ramps * Purpose: Perform residual calibration and check for stability of polarizer and ramp filters. * Description: This proposal does not duplicate the existing ramp wavelength calibration or polarization calibration (Cycle 4 programs 5574 and 6140). Instead, it provides a full polarization calibration for a filter that was not supported in Cycle 4 (F30OW), a check for polarization stability, and a throughput calibration for the linear ramp filters, by scanning the spectrophotometric standard along the ramps. * Accuracy: The proposal should support polarimetry at the 3 percent level and measure the ramp throughput at the 2 percent level. 6195: Flat field Check * Purpose: Check quality of flat fields and estimate errors in them. * Description: The crowded omega Cen field is positioned with a bright star at the center of each CCD in turn. 40 second images are taken through each of the photometric filters (F336W, F439W, F555W, F675W, F814W) and as many supplementary filters (from F450W, F606W, F702W and F547M) as can be fitted in. If data volume is a problem, single chip readout is acceptable, but should be avoided as much as possible. This is based on Cycle 4 programs 5659 and 5646. * Accuracy: Overall discrepancies between the existing flat fields and the results of this test should be 1-2 percent rms. Part of the point of the test is to measure this accuracy. * Products: After pipeline processing, each image will be reduced by aperture photometry to measure the RMS errors in the flat fields. The RMS error will, be determined by the additional noise in the independent measurements over the expected variance of less than 1 percent from photon statistics. The single bright star at the center of each chip independently estimates the chip to chip normalization error. 6250: Internal Monitor * Purpose: Check for short term stability of instrument. * Description: The routine internal monitor, to be run twice weekly for WFPC2 during Cycle 5, obtains two bias frames at each gain, two INTFLATs with the F555W filter at each gain, and two Kelsall spot images with exposure times optimized for the WF and PC, respectively. It is identical to the Cycle 4 program 5560. * Accuracy: This monitor is not involved in generating quantitative calibration information. * Products: The test provides a biweekly monitor of the integrity of the CCD camera chain electronics both at gain 7 and 15, a test for quantum efficiency hysteresis in the CCDS, and an internal check on the alignment of the WFPC2 optical chain. Stiavelli and Baggett, WFPC2 ISR 96-01. 8.10 Cycle 6 Calibration Plan The Cycle 6 calibration plan is similar to that for Cycle 5, and is summarized in Table 8.3. Important differences include the addition of programs to check the astrometric calibration (6941), more detailed checking of the camera electronics (6942), and measurements of narrow band filter throughputs (6943). Also, it is expected that there will be reduced usage of the calibration channel VISLAMP, so as to prolong the lamp lifetime. More detailed program descriptions are given below. Table 8.3: Summary of Cycle 6 Calibration Plan. External ID Proposal Title Schedule Results Accuracy Time Notes (percent) (orbits) ------------------------------------------------------------------------------ Routine Monitoring Programs 6902 Photometric 2x per 4 weeks SYNPHOT 2 26 Monitor 6903 Decontamination 1 x per 4 weeks CDBS n/a 0 darks, internals 6904 Darks Weekly CDBS 1 e/hr 0 WWW hot pixel lists 6905 Internal Monitor Weekly CDBS 0.8 c 0 6906 Visflat Monitor 2 x per 4 weeks ISR 0.3 0 (monitor lamp health) 6907 Intflat Monitor 1 x per 4 weeks ISR 0.3 0 6908 UV Flat Field 2 in Cyc. 6 ISR 2-8 0 Monitor 6909 Earth Flats Continuous CDBS 0.3 0 Special Calibration Programs 6934 Photometric l x in Cyc. 6 SYNPHOT 2 6 Add 2 new Zeropoint standards 6935 Photometric Trans. 2 x in Cyc. 6 ISR 2-5 9 Three targets 6936 UV Throughput & 2 x in Cyc. 6 SYNPHOT 3-10 12 Include Ly alpha BD+75D325 6937 CTE Calibration 1 x in Cyc. 6 ISR 1 2 6938 PSF 1 x in Cyc. 6 CDBS 10 7 Characterization 6939 Linear Ramp 1 x in Cyc. 6 CDBS 3 4 Filters 6940 Polarizers 1 x in Cyc. 6 CDBS 3 4 6941 Astrometry 1 x in Cyc. 6 STSDAS 0.01" 4 Verification 6942 Camera Elect. 1 x in Cyc. 6 ISR 0.5 1 Verification 6943 Narrow Band SNAP SYNPHOT 3 8 SNAP Throughput ------------------------------------------------------------------------------ TOTALS 83 ------------------------------------------------------------------------------ 6902: Photometric Monitor * Purpose: Monthly external check of instrumental stability. * Description: The standard star GRW+7OD5824 is observed before and after each decontamination (thus twice in a four-week period). Each observation consists of three sequences: (1) F170W in all four chips to monitor contamination in the far UV; (2) F439W, F555W, F814W on the PC to monitor focus; and (3) F160W, F185W, F218W, F255W, F300W, F336W, F439W, F555W, F675W, F814W in a different chip each month. Some filters may be cut because of lack of time (FI85W cut first, then F30OW, then F675W, then F218W). This proposal is based largely on Cycle 5 program 6184; focus monitoring in F439W and F814W is added at the expense of some UV filters. * Accuracy: Overall discrepancies between the results of this test need to be measured to better than 2 percent and are expected to be less than 1 percent rms. The point of the test is to measure this variation. * Products: Documents produced are Instrument Handbook, TIPS, (sensitivity trends). Updates in UV sensitivity variation used in SYNPHOT are provided. 6903: Decontamination * Purpose: UV blocking contaminants are removed by warming the CCDS. * Description: The decon itself is implemented via use of the DECON mode, in which the TECs are turned off and the CCD and heatpipe heaters are turned on to warm the detectors and window surfaces. Keeping WFPC2 warm for ~6 hours has been shown in previous Cycles to be sufficient to remove the contaminants and anneal many hot pixels; continuation of 6-hour decons is anticipated for Cycle 6. The internal observations taken before and after each decontamination consist of: 4 biases (2 at each gain setting), 4 INTFLATs (2 at each gain setting), 2 kspots (both at gain 15, one short and one long exposure, optimized for PC and WF), and finally, 5 darks (gain 7, CLOCKS=NO). To minimize time-dependent effects, each set of internals will be grouped within 2 days and performed no more than 1 day before the DECON and no later than 12 hours after the DECON. To protect against residual images in the darks (which results in the irretrievable loss of the critical pre-DECON hotpixel status), the darks will be executed NON-INT and requested to be done at least 30 minutes after any WFPC2 activity. Special Requirements: This requires scheduling at 4 week intervals. It prevents WFPC2 from being used for several hours, although other instruments can be used most of that time. Dark frames taken before decontaminations need to be protected from possible residual images from overexposed sources. * Accuracy: This proposal is mainly designed to maintain the health of the instrument. Biases, darks and other internals taken with this proposal are used in generating appropriate reference files. * Products: They are obtained from use of darks, biases and other internals (see Proposals 6904 and 6905). 6904: Darks * Purpose: Measure dark current on individual pixels and identify hot pixels at frequent intervals. * Description: Every week, five 1800s exposures are taken with the shutter closed. The length of the exposures is chosen to fit nicely within an occultation period. The weekly frequency is required because of the high formation rate of new hot pixels (about 70/CCD/day). Five darks a week are required for cosmic ray rejection, to counterbalance losses due to residual images, and to improve the noise of individual measurements. Even with these measures, some weeks no usable darks will be available because of residual images. Normally this results only in a longer-than-usual gap in the hot pixel lists, but in a decontamination week, information on pixels that became hot and then annealed would be lost irretrievably. For this reason, pre-DECON darks are to be executed NON-INT and at least 30 minutes after any WFPC2 activity (see Proposal 6903). Normal darks do not need to be protected in this fashion. * Accuracy: Superdarks should be accurate to better than 1 e/hour and are expected to reach errors of about 0.05 e/hour (single-pixel rms). Systematic errors due to dark glow (a spatially and temporally variable component of dark signal) and hot pixels may exceed these limits significantly. * Products: Weekly dark frames are delivered to CDBS and monthly tables of hot pixels are posted on the Web. 6905: Internal Monitor * Purpose: Verification of short-term instrument stability for both gain settings. * Description: The internal observations will consist of 8 biases (4 at each gain, and 4 INTFLATs (2 at each gain). The entire set should be run once per week (except for DECON weeks) on a non-interference basis. This proposal is similar to the Cycle 5 Internal Monitor (6250), except that the K-spot images have been removed (these are being taken with the DECON Proposal). The execution frequency during Cycle 6 has also been reduced, from twice a week to once a week, although the total number of biases has been increased to continue to provide an adequate number of frames for pipeline reference file generation. * Accuracy: Approximately 40 bias frames will be used for each pipeline reference file; accuracy is required to be better than 1.5 e/pixel, and is expected to be 0.8 e/pixel. * Products: Superbiases are delivered every few months to CDBS. TIPS reports are made on possible buildup of contaminants on the CCD windows (worms) as well as gain ratio stability, based on INTFLATS. ISRs are prepared if significant changes occur. 6906: Visflat Monitor * Purpose: Monitor the stability of the camera and filter responses via the VISFLAT channel. * Description: Twice a month, internal flat fields (VISFLATS) will be obtained using the visible calibration channel lamp with the photometric filter set plus a couple of narrow-band filters. The images will be used to monitor WFPC2's flat field response as well as to build a high S/N flat field database, which will provide information on the pixel-to-pixel response in the cameras and any possible long-term contamination-induced changes. The LRF (FR533N) exposures, one at each gain, taken after DECON will provide a monitor of the ADC's performance. Histograms generated from the ramp filter flats will be used to trace the ADC transfer curve. ON HOLD: In addition to the monitor observations, an initial filter-sweep is done to obtain VISFLATs in all visible filters. These will be compared to the Cycle 5 filter-sweep data to verify that none of the filters are developing any problems, and to provide a check of the calibration channel's long-term stability. * Special Requirements: Uses the VISFLAT calibration channel, whose Welch-Allyn bulb is apparently wearing out. (A back-up exists for the Welch-Allyn bulb.) The Cycle 6 proposal has been redesigned to limit the number of ON/OFF cycles placed on this channel to a level believed safe over 10-15 years. The sweep part of the proposal, which puts the heaviest usage on the lamp, is on hold, pending verification of the lamp health from the short monthly executions. The INTFLAT Monitor (Proposal 6907) can obtain similar information if necessary. * Accuracy: The VISFLAT response is stable to about 0.3 percent, both in overall level (lamp degradation aside) and in spatial variations. The point of this proposal is to verify this stability on a regular basis and to monitor the lamp degradation. * Products: ISR and TIPS reports will be prepared. 6907: Intflat Monitor * Purpose: Provide backup database of INTFLATS in case VISFLAT channel fails. * Description: This proposal consists of two parts: 1) an INTFLAT filter sweep and, 2) a series of exposure to test the linearity of the camera. 1) The sweep is a complete set of internal flats cycling through both shutter blades and both gains. Signal-to-noise per pixel is estimated to be similar to the VISFLATs (0.6 percent) but the spatial and wavelength variations in the illumination pattern are much larger. However, the INTFLATs will provide a baseline comparison of INTFLAT vs. VISFLAT, in the event of a calibration channel system failure and temporal variations in the flat fields at the 1 percent level. In addition, these images will provide a good measurement (better than 1 percent) of the stability of the gain ratios. 2) The linearity test portion is aimed at obtaining a series of INTFLAT with both gains and both shutters. Since the INTFLATs have significant spatial structure, any non-linearity would appear as a non-uniform ratio of INTFLATs with different exposure times. A set of exposures is also taken with gain 7, shutter B, and CLOCKS=YES. * Accuracy: The signal-to-noise per pixel is similar to that obtained in the VISFLAT program (0.6 percent) but there are much larger spatial and wavelength variations in the illumination pattern. As a result, this dataset will not form any part of the pipeline calibration. This baseline is necessary in case the VISFLAT channel fails and there are temporal variations in the camera flatfields at the 1 percent level. * Products: TIPS reports and ISRs will be prepared if any significant variations are observed. 6908: UV Flat Field Monitor * Purpose: Monitor the stability of UV flat field. * Description: UV flat fields will be obtained with the calibration channel's ultraviolet lamp (UVFLAT) using the UV filter set (FI22M, F170W, F16OBW, F185W, and F336W). The UV flats will be used to monitor UV flat field stability and the stability of the Woods filter (FI60BW) by using F170W as the control. The F336W ratio of VISFLAT (Cycle 6 proposal 6906)/UVFLAT ratio will provide a diagnostic of the UV flat field degradation and tie the UVFLAT and VISFLAT flat field patterns together. Two supplemental dark frames must be obtained immediately after each use of the lamp, in order to check for possible afterimages. * Special Requirements: This uses the limited life UV lamp. In order to prevent excessive degradation of the lamp, the SU duration for each UVFLAT visit should be kept the same as the durations used during Cycle 5 (proposal 6191); the lamp should not remain on for periods of time longer than those used in Cycle 5. To be executed once before and once after the refurbishment mission, shortly after a decontamination. * Accuracy: About 2-8 percent pixel-to-pixel are expected (depending on filter). * Products: New UV flat fields are made if any changes are detected. 6909: Earth Flats * Purpose: Monitor flat field stability. * Description: As in Cycle 5 program 6187, 4 sets of 200 Earth streak-flats are taken to construct high quality narrow-band flat fields with the filters F160BW, F375N, F502N, F656N and F953N. Of these 200 perhaps 50 will be at a suitable exposure level for de-streaking. The resulting Earth super flats map the OTA illumination pattern and will be combined with SLTV data (and calibration channel data in case of variation) for the WFPC2 filter set to generate a set of superflats capable of removing both the OTA illumination and pixel-to-pixel variations in flatfield. The Cycle 4 plan is being largely repeated except: (1) UV filters are dropped because measurement is generally only of the redleak; (2) F160BW is retained in order to check for developing pinholes; and (3) Crossed filters used as neutral densities are eliminated (illumination pattern is wrong). Specific observations for the Methane Quad filters will also be included. * Accuracy: The single-pixel noise expected in the flat field is 0.3 percent. * Products: New flat fields to CDBS if changes detected. 6934: Photometric Zeropoint * Purpose: Verify synthetic zeropoint of WFPC2 filters. * Description: Standard stars are observed through all filters longward of F336W (inclusive) with the limit of one orbit per target. Targets include: the spectrophotometric standard GRW+70D5824 in PC and WF3; two stars chosen as standards by other instruments on HST; and two standard star fields, containing 3-4 stars each, commonly used in ground-based photometry. Observations of GRW+70D5824 will be directly comparable to the corresponding observations for cycles 4 and 5 (programs 5572 and 6179) and will verify the stability of the filters as well as improve the accuracy of the calibration. The other standards are observed to provide cross-instrument calibration and in order to increase the range of colors used for photometric verification. This proposal will help all observers who want to do quantitative photometry at the 2 percent level. * Special Requirements: Specific orientations will be required for the two fields of standards in order to fit the maximum number of stars. All observations should be executed within a week after decontamination. * Accuracy: 2 percent required, 1 percent expected for our main spectrophotometric standard GRW+7OD5824. Expected accuracy for the other standards is between 2 percent and 5 percent, depending on spectral type and filter; most of the error derives from limited knowledge of the transformations between ground-based and WFPC2 photometric systems. * Products: TIPS reports, SYNPHOT updates if necessary, and ISRs will be prepared. 6935: Photometric Transformation * Purpose: (1) Update photometric transformations to Johnson-Cousins system and Str6mgren system; (2) Deten-nine spatial dependence of contamination; (3) Check the astrometric solution using M67; (4) Spot check of gain=7 vs. gain= 15 ratios; 5) Spot check short vs. long exposure zeropoints. * Description: Three photometric standard star fields in NGC 5139 (omega Cen; metal rich), NGC 2682 (M67; metal poor), and NGC 2100 (young cluster) are observed before and after a decontamination. Four different filter sets are used: (1) The five filters generally used to match the Johnson-Cousins system (F336W, F439W, F555W, F675W, F814W); (2) The wide-band equivalents for the Johnson-Cousins system (F30OW, F380W, F450W, F606W, F702W); (3) The Str6mgren equivalents (F336W, F410M, F467M, F547M); and (4) Two filters farther toward the UV (F255W, F170W), so that contamination over the full field of view can be measured. F255W is not used for the reddest cluster (NGC 5139). F17OW is only used for the bluest cluster (NGC 2100). For the brighter clusters (NGC 2682 and NGC 2100) long and short exposures are taken in the UBVRI equivalents both to extend the dynamic range and to check for differences in photometric zeropoints. A spot check is included to compare gain=7 and gain=15 is also included for NGC 5139 and NGC 2682. * Special Requirements: The first visit for each target must be taken within 3 days after a decontamination. The second visit, including only the UV filters, must be taken more than 25 days after the first visit, but before the next decontamination. * Accuracy: The photometric transformations should be accurate to 2-5 percent The stability of these transformations will be measured to the 1 percent level. The astrometry should be good to 0.1" (absolute) and 0.05" (relative). * Products: ISR and Instrument Handbook. It will also be part of a planned paper on the possibility to do 1 percent photometry. 6936: UV Throughput and Lyman-alpha Verification * Purpose: Verify throughput for all UV filters, including Lyman-(X test to monitor possible contamination on pick-off mirror. * Description: Spectrophotometric standards are observed shortly before and after a DECON through all the UV filters in each chip and through F160BW crossed with F130LP, F185LP, and F165LP to determine the wavelength dependence of the throughput across the bandpass (for color terms). This proposal is based on the Cycle 5 UV Throughput proposal (6186) but includes also the standard BD+75D325 used in Cycle 4 (proposal 5778) to establish the Lyman-alpha throughput calibration. * Special Requirements: Timing requirements with respect to decontaminations. * Accuracy: The UV throughput will be measured to better than 31 percent. Accuracy in Lyman-alpha throughput is expected to be between 5 and 10 percent, because of the residual uncertainty of the red leak correction after observations with crossed F122M and F13OLP. * Products: TIPS reports, SYNPHOT updates, and ISRs will be prepared if necessary. 6937: CTE Calibration * Purpose: (1) Test whether the exposure-time dependence of the photometric calibration is due to CTE (long vs. short exposure problem); (2) refine flux and background-level dependent aperture corrections. * Description: The globular cluster NGC 2419 will be observed through F555W with a combination of exposure times (between 5 and 1400 s) and preflash levels (0 to 500 e-). Analysis of cycle 5 CTE calibration (proposal 6192) suggests that magnitude errors due to charge transfer effects are greatly reduced (if not entirely eliminated) at background levels of 160 e- or greater. CTE effects have been proposed as the solution to reported differences in the magnitudes of stars measured on frames with short exposures vs. long exposures. If CTE is the cause, then the differences should disappear with preflash. We will re-observe the cluster NGC 2419 with and without preflash to test this hypothesis. This dataset will also provide a large number of stars with which to refine existing measurements of the effects of CTE on the wings of the PSF. A range of preflash levels will be explored at the 60 sec exposure time. * Special Requirements: Observations should be made at the same position and roll angle as the previous NGC 2419 LONG exposures (proposal GO-5481). * Accuracy: The reported short vs. long effect is ~0.05 mag. We wish to reduce this to less than 0.01 mag. * Products: An ISR will be prepared. If appropriate, a special task to correct the CTE effect will be generated. 6938: PSF Characterization * Purpose: Provide a sub-sampled PSF over the full field to allow PSF fitting photometry, test PSF subtraction as well as dithering techniques (c.f. effects of the OTA breathing and CCD gain). * Description: Measure PSF over full field in photometric filters in order to update the TIM and TINYTIM models and to allow accurate empirical PSFs to be derived for PSF fitting photometry. These observations will also be useful in order to test PSF subtraction and dithering techniques at various locations on the CCD chips. With one orbit per photometric filter, a spatial scan is performed over a 4X4 grid on the CCD. The step size is 0.025 arcseconds; this gives a critically sampled PSF over most of the visible range. This program uses the same specially chosen field in omega Cen as the Cycle 5 proposal 6193, but with a few arcsec shift in order to map the PSF variation better. The standard 'photometric' filters are used. Two additional orbits are used to explore the effects of OTA breathing and CCD gain onto dithering and PSF subtraction techniques. Data volume will be a problem, so special tape recorder management will be called for. The proposal also allows a check for sub-pixel phase effects on the integrated photometry. * Special Requirements: This needs the same pointing and orientation as Cycle 5 observations for proposal 6193, thus should be scheduled within a similar time frame. * Accuracy: It provides measurement of pixel phase effect on photometry (sub pixel QE variations exist). The chosen field will have tens of well exposed stars in each chip. Each star will be measured 16 times per filter at different pixel phases. The proposal therefore provides, in principle, a high signal-to-noise, critically sampled PSF. This would leave PSF fitting photometrists in a much better position than now, where pixel undersampling clearly limits the results. The result will be largely limited by breathing variations in focus. It is hard to judge the PSF accuracy that will result. If breathing is less than 5 microns peak-to-peak, the resulting PSFs should be good to about 10 percent in each pixel. Breathing effects will be investigated (1 additional orbit) as well as the gain dependence (I additional orbit). PSF fitting results using this calibration would of course be much more accurate. In addition, the test gives a direct measurement of sub-pixel phase effects on photometry, which should be measured to better than 1 percent. * Products: A PSF library (CDBS) will be assembled, and an ISR will be issued as needed. 6939: Linear Ramp Filters * Purpose: Verify throughput calibration for Linear Ramp Filters at selected wavelengths. * Description: Throughput calibration is obtained by observing the spectrophotometric standard GRW+70D5824 at several filter rotations and wavelengths. This completes the program carried out in Cycle 5, in which some wavelength ranges and rotations could not be covered. * Accuracy: Throughput accuracy should be verified to better than 3 percent; 1-2 percent can be achieved. * Products: SYNPHOT throughput tables will be updated if necessary. 6940: Polarizers * Purpose: Verify stability of polarization calibration. * Description: The goal of this proposal is to check for any changes in the polarization calibration since Cycle 5. Observations are made in F555W+POLQ of both polarized and unpolarized stars, in addition to VISFLATS. Data are taken in all four quads of the polarizer, as well as in three rotated positions of the POLQ. * Special Requirements: Requires specific orientations. * Accuracy: 3 percent. * Products: Update throughput tables if necessary 6941: Astrometry Verification * Purpose: Verify accuracy and stability of geometric transformation and relative astrometry solution. * Description: A very rich star field in co Cen will be observed in five different positions with relative shifts of 40' in each coordinate. Positions of more than 2000 stars per chip will be compared between paintings using the three different astrometric solutions provided by Gilmozzi, Holtzman and Trauger, in order to verify and refine their accuracy. (Differences of up to 1 PC pixel exist in some regions of the field of view.) A very densely populated field is chosen in order to achieve better coverage of the field of view, even if at the expense of the accuracy of individual position measurements. Observations will be carried out in three filters, F555W, F300W, and F814W, to provide a verification and/or correction of the wavelength dependence of the solution. The F555W observation is repeated with smaller shifts of 15" to ensure a better coverage of the PC. * Accuracy: We expect better than 0.01" and require better than 0.05" (full field of view). * Products: Improvements will be noted in METRIC and in the aperture reference file if required. An ISR will be issued as needed. 6942: Camera Electronics Verification * Purpose: Verify several aspects of the WFPC2 camera electronics: linearity, gain ratios, effect of CLOCKS, and effect of CTE on extended sources. * Description: Observing a very extended nonuniform target represented by the giant elliptical galaxy NGC 4472. The linearity test is carried out by taking exposures of NGC 4472, centered in WWALL, with a variety of exposure times. Since the galaxy is non-uniform, the ratio of these exposures is directly related to the camera linearity. The exposures will be taken with GAIN=7. However, one exposure will also be taken with GAIN=15. Additional exposures will be taken with CLOCKS=YES and with a preflash. These observations complement those of the internal calibration proposal 6907. The two major advantages of these observations compared to the 6907 ones are the possibility of studying the effect of preflash (since the light distribution of the preflash is different from that of NGC 4472) and the possibility of measuring an absolute response curve, since NGC 4472, unlike the INTFLAT lamp, does not have variations in luminosity. NGC 4472 has been chosen as target galaxy because it is large enough to produce significant signal in all chips and bright enough to allow us to explore the highest counts without excessive integration times. * Accuracy: We expect 0.5 percent for linearity and CTE, 0.1-0.2 percent for gain ratios and CLOCKS. We require less than 1 percent on each item. * Products: ISRs and TIPS reports will be prepared as needed. 6943: Throughput Verification for Narrow Band Filters * Purpose: Direct verification of throughput of narrow band filters through observations of emission line objects. * Description: The current throughput calibration of narrow-band filters is based on filter profiles from data obtained before launch and on observations of continuum sources. This program will verify the accuracy of the calibration, and indirectly the stability of the filters, by observing eight planetary nebulae with strong lines and well-established ground-based spectra. The observations can be executed in SNAPSHOT mode since they will be short and none is specifically required. Some planetary nebulae with existing Cycle 4 and 5 observations will be included for stability verification. * Accuracy: We expect 2 percent and require 3 percent. * Products: SYNPHOT tables will be updated and an ISR issued if required. ------------------------------------------------------------------------------ CHAPTER 9: References In This Chapter... References Instrument Science Reports 9.1 References * Bessell, M. S. (I 990), UB VRI Passbands, RA.S.P. 102, 118 1. * Biretta, J. A, and Sparks, W. B. (1995), WFPC2 Polarization Observations: Strategies, Apertures, and Calibration Plans, WFPC2 ISR 95-01. * Biretta, J. A., (1995) "WFPC2 Flat Field Calibration," in Calibrating Hubble Space Telescope: Post Servicing Mission, eds. A. Koratkar and C. Leitherer, p. 257. * Blouke, D. P., (1991), eds. Janesick, J. and Elliot, T., in P.A.S.P., 8, 153. * Burrows, C. et al., (1991) The Imaging Performance of the Hubble Space Telescope, Ap. J. Lett. 369, L21. * Clarke, J.T. and the WFPC2 IDT (1992), "White Paper for WFPC2 Far-Ultraviolet Science."^1 * Dithering: Relationship Between POS TARGs and CCD Rows/Columns.^1 * Evans, R. E. 1992, JPL Memorandum, DFM #2031. * Gilliland, R. L., (1994) Ap.J. 435, L63. * Groth, E. and Shaya, E., (1991), in Wide Field/Planetary Camera Final Orbital/Science Verification Report, S., M., Faber, Ed. * Griffiths, R. Ewald, S. and MacKenty, J. W., (1989) in CCDs in Astronomy, P.A.S.P. 8, 231, ed. G. Jacoby. * Gunn, J. E., and Oke, J. B. (I 983), Ap.J. 266, 713. * Harris, H. C., et al., (1991) Photometric Calibration of the HST WF/PC-1 Camera: I. Ground Based Observations of Standard Stars, A.J. 101, 677. * HST Data Handbook, C. Leitherer, ed., (Version 2.0, December 1995).^1 * HST Phase 2 Proposal Instruction, (Version 8.0, 15 December 1995).^1 * Hasan, H. and Bely, P., (1993) Restoration of HST Images and Spectra II, p. 1 57. * Hasan, H. and Burrows, C. J., (1993) "Calibrations of the Hubble Space Telescope Optical Telescope Assembly" in Calibrating Hubble Space Telescope, eds. C. Blades and S. Osmer, p. 395. * Harris, et al., (1991) A.J. 101, 677. * Holtzman, J., et al., (1995a), The Performance and Calibration of WFPC2 on the Hubble Space Telescope, P.A.S.R, 107, 156. * Holtzman, J., et al., (1995b), The Photometric Performance and Calibration of WFPC, P.A.S.P., 107,1065. * Interface Control Document (ICD) 19, "PODPS to STSDAS". * Interface Control Document (ICD) 47, "PODPS to CDBS". * Janesick, J., Elliot, T., Bredthauer, R., Cover, J., Schaefer, R. and Varian, R., (1989) in Optical Sensors and Electronic Photography, SPIE Proc. 1071. * Janesick, J., Elliot, T., Blouke, M. and Corrie, B., (1989b), in Optical Sensors and Electronic Photography, SPIE Proc. 1071, 115, ed. B. Pophal. * Jordan, Deltom, and Oates, (1993) Greenwich Observatory Newsletter, Sept. 1993. * Koornneef, J. et al., (1983) "Synthetic Photometry and the Calibration of the Hubble Space Telescope," in Highlights of Astronomy, 7, 833, ed. J.-P. Swings. * Krist, J. and Burrows, C., (1995) Applied Optics. * Krist, J. E., (1995) "WFPC2 Ghosts, Scatter, and PSF Field Dependence," in Calibrating Hubble Space Telescope: Post Servicing Mission, eds. A. Koratkar and C. Leitherer, p. 311. * Lauer, T. (1989), The Reduction of WFIPC Camera Images, P.A.S.R 101, 445. * Lauer, T., (1991) in Wide FieldIPlanetary Camera Final Orbital/Science Verification Report, ed. S. M. Faber. * Oke, J.B. and Gunn, J., (1983), Ap. J. 266, 713. * STSDAS Calibration Guide (November 1991).^1 * STDAS Users' Guide. * Schneider, D. P, Gunn, J. E., and Hoessel, J. G. (1983), Ap.J. 264, 337. * Trauger, J. T., (1989), in CCDs in Astronomy, P.A.S.P. 8, 217, ed. G. Jacoby. * Trauger, J. T., ed., (1993). The WFPC2 Science Calibration Report, Pre-launch Version 1.2. [IDT calibration report] * WF/PC-1 IDT OV/SV Report. * WFPC2 World Wide Web page at address: http://www.stsci.edu/ftp/instrument_news/WFPC2/wfpc2_top.html * Westphal, J., et al. (1982), The Wide Field/Planetary Camera in The Space Telescope Observatory, IAU 18th General Assembly, Patras, NASA CP-2244. * Woodgate, B. E., (1989) in CCDs in Astronomy, P.A.S.P. 8, 237, ed. G. Jacoby. 1. Documents may be requested by e-mail from help@stsci.edu Instrument Science Reports These documents may be requested by e-mail from help@stsci.edu. * 96-02: Baggett, et al., Contamination Correction in SYNPHOTfor WFPC2 and WF/PC-1. * 96-01: Stiavelli and Baggett, 01/96, Internal Flat Field Monitoring. * 95-07: Baggett, Casertano, and Biretta, 12/95, WFPC2 Cycle 4 Calibration Summary. * 95-06: Biretta, Ritchie, and Rudloff, 08/95, A Field Guide to WFPC2 Image Anomalies. * 95-05: Biretta, Ritchie, Baggett, MacKenty 04/96, Wavelength/Aperture Calibration of the WFPC2 Linear Ramp Filters. * 95-04: Whitmore and Heyer, 07/95, Demonstration Analysis Script for Performing Aperture Photometry. * 95-03: Whitmore and Wiggs, 07/95, Charge Transfer Traps in the WFPC2. * 95-02: Gilmozzi, R. et al., 06/95, The Geometric Distortion of the WFPC2 Cameras. * 95-01: Biretta and Sparks, 02/95, WFPC2 Polarization Observations: Strategies, Apertures, and Calibration Plans. * 094-03: Burrows, C., 12/94, WFPC2 Pipeline Calibration. * 94-01: Krist, J. and Burrows, C., 10/94, Large Angle Scattering in WFPC2 and Horizontal "Smearing" Correction. * 93-01: Clampin, M., 3/93, Polarizer Quad Nomenclature. * 92-06: Clampin, M., 12/92, WFPC2 CCDS. * 92-05: Burrows, C., 11/92WFPC2 AFM and POMM Actuation Algorithm. * 92-04: Trauger, J.T. et al., 7/92, Science with the Second Wide Field and Planetary Camera. * 92-03: Trauger, J.T. and Brown, D.I., 10/92, WFPC2 Science Observation and Engineering Modes. * 92-02: Clampin, M., 9/92, System Level Contamination Issues for WFPC2 and COSTAR.