The Solar Wind Around Pluto (SWAP) Instrument Aboard New Horizons

D. McComas (1,*), F. Allegrini (1), F. Bagenal (2), P. Casey (1),
P. Delamere (2),
D. Demkee (1), G. Dunn (1), H. Elliott (1), J. Hanley (1), K. Johnson (1),
J. Langle (3,1), G. Miller (1), S. Pope (1), 
M. Reno (1), B. Rodriguez (1), N. Schwadron (4,1), P. Valek (1),
S. Weidner (1)

1 Southwest Research Institute, San Antonio, Texas 78228-0510, U.S.A.
2 University of Colorado, Boulder, Colorado 80309, U.S.A.
3 Midwest Research Institute, Kansas City Missouri 64110, U.S.A.
4 Boston University, Boston Massachusetts 02215, U.S.A.
( *author for correspondence, e-mail: DMcComas@swri.edu)

Abstract. The Solar Wind Around Pluto (SWAP) instrument on New
Horizons will measure the interaction between the solar wind and ions
created by atmospheric loss from Pluto. These measurements provide a
characterization of the total loss rate and allow us to examine the
complex plasma interactions at Pluto for the first time. Constrained
to fit within minimal resources, SWAP is optimized to make plasma-ion
measurements at all rotation angles as the New Horizons spacecraft
scans to image Pluto and Charon during the flyby. In order to meet
these unique requirements, we combined a cylindrically symmetric
retarding potential analyzer (RPA) with small deflectors, a top-hat
analyzer, and a redundant/coincidence detection scheme. This
configuration allows for highly sensitive measurements and a
controllable energy passband at all scan angles of the spacecraft.

Keywords: solar wind, Pluto plasma interaction, pickup ions

1.0 Introduction

The New Horizons mission [Stern et al., 2007] will make the first
up-close and detailed observations of Pluto and its moons. These
observations include measurements of the solar wind interaction with
Pluto. The Solar Wind Around Pluto (SWAP) instrument is designed to
measure the tenuous solar wind out at ~32 AU and its interaction with
Pluto. SWAP directly addresses the Group 1 science objective for the
New Horizons mission to "characterize the neutral atmosphere of Pluto
and its escape rate."  In addition, SWAP makes measurements critical
for the Group 2 objective of characterizing Pluto's ionosphere and
solar wind interaction.  Finally, SWAP makes observations relevant to
two Group 3 science objectives for New Horizons, both in support of
characterizing the energetic particle environment of Pluto and Charon
and in searching for magnetic fields of Pluto and Charon.

Pluto's thick atmosphere escapes the planet's weak gravity and streams
away as neutral particles. The escaping neutral atoms are ionized by
solar photons. As the solar wind approaches Pluto it picks up these
ions and electrons and carries them downstream. This pickup process
slows the solar wind, and the type of interaction varies greatly
depending on the atmospheric escape rate, all the way from Comet-like
for larger escape rates to Venus-like for low escape rates. SWAP
measurements should provide the best estimate of the overall
atmospheric escape rate at Pluto and allow the first detailed
examination of its plasma interactions with the solar wind.

The SWAP observations are extremely challenging because the solar wind
flux, which falls off roughly as the square with heliocentric
distance, is approximately three orders of magnitude lower at Pluto
compared to typical solar wind fluxes observed near Earth's orbit. In
addition, because the solar wind continues to cool as it propagates
out through the heliosphere, the solar wind beam becomes narrow in
both angle and energy.

The SWAP design was strongly driven by three constraints: 1) very low
use of spacecraft resources (mass, power, telemetry, etc.) as this is
a non-core instrument on a relatively small planetary mission; 2) very
high sensitivity to measure the solar wind and its interaction with
Pluto out at ~32 AU, where the density is down by a factor of ~1000
compared to 1 AU; and 3) the need to make observations over a very
large range of angles as the spacecraft constantly repoints its
body-mounted cameras during the flyby. Given these
not-entirely-consistent design drivers, we developed an entirely new
design that combines elements of several different previous plasma
instruments. Together these components comprise the SWAP instrument,
which will measure the speed, density, and temperature of the distant
solar wind and its interaction with Pluto.

2.0 Scientific Background and Objectives

Only partly in jest, Dessler and Russell (1980) suggested that Pluto
might act like a colossal comet. The 1988 stellar occultation showed
that Pluto's tenuous atmosphere could indeed be escaping (Hubbard et
al., 1988; Elliot et al., 1989). Applying basic cometary theories to
Pluto, Bagenal and McNutt (1989) showed that photoionization of
escaping neutral molecules from Pluto's atmosphere could significantly
alter the solar wind flow around Pluto for sufficiently large escape
rates. For large atmospheric escape rates, the interaction may be best
described as "comet-like," with significant mass-loading over an
extensive region; for small escape rates the interaction is probably
confined to a much smaller region, creating a more "Venus-like"
interaction (Luhmann et al., 1991), where electrical currents in the
gravitationally bound ionosphere deflect the solar wind flow. Figure 1
compares these two types of interactions schematically. At aphelion
(50 AU), should Pluto's atmosphere completely collapse and freeze onto
the surface, then the interaction becomes "Moon-like" with the solar
wind suffering minimal deflection and directly bombarding the bare,
icy dayside surface. Not having a detectable atmosphere, Charon almost
certainly has such a "Moon-like" interaction, remaining primarily in
the solar wind if Pluto's interaction is weak but becoming totally
engulfed if Pluto's interaction is strongly "comet-like" and extends
beyond Charon's orbit at 17 RP (Pluto radii, ~1150 km). For a review
of early studies of the solar wind interaction with Pluto see Bagenal
et al. (1997), which also reviews the implications of the unlikely
possibility of Pluto having an intrinsic magnetic field.

Figure 1. Solar wind interaction with Pluto for (a) low atmospheric
escape rate and (b) high atmospheric escape rate. The dots indicate
ions produced by the photoionization of Pluto's escaping
atmosphere. These pick-up ions will move perpendicular to both the
magnetic field (B) and solar wind (upwards in the top two
diagrams). Note that while the IMF tends to be close to tangential to
Sun direction, the sign of the direction varies on timescales of
days. The asymmetries of the interaction will flip as the magnetic
field changes direction. The lower panels show a plane parallel to the
ecliptic plane though at the time of the New Horizons flyby Pluto will
be well above the ecliptic plane.

The solar wind is supersonic (actually superfast-mode magnetosonic) so
that when it impinges on a magnetic obstacle (such as the
magnetosphere of the Earth or other planet) an upstream bow shock must
form to slow and deflect the supersonic plasma. The weak
interplanetary magnetic field (IMF) at 30 AU (see Table I for typical
solar wind properties near Pluto) and heavy ions formed by
photoionizing the heavy molecules of Pluto's escaping atmosphere have
very large gyroradi (~500 RP). The net results of these non-fluid or
kinetic effects are to make the bow shock a thick transition region
and to make the shape of the interaction region asymmetric where the
direction of asymmetry is governed by the direction of the IMF. Recent
simulations of the solar wind interaction with a strongly escaping
atmosphere have necessarily been 3D and have either taken a
multi-fluid approach (solar wind proton fluid and pick-up ion fluid)
or a hybrid approach (electron fluid, ion particles) (Harnett et
al. 2005; Delamere and Bagenal, 2004). 

Table I. Typical Solar wind Interaction Properties at Pluto (at 30
AU), taken from Bagenal et al. (1997) using Rp=1150 km. 

Below we discuss how the current understanding of Pluto's atmosphere
leads us to expect a more comet-like interaction at the time of the
New Horizons flyby in 2015. Through measurements of bulk properties of
the solar wind (flow, density, temperature) as well as the energy
distribution of solar wind and pick-up ions, the SWAP instrument will
not only characterize the solar wind interaction with Pluto but will
also allow us to determine the global rate of atmospheric escape.

In the comet-like scenario, variations on the scale of the interaction
region can be substantial over periods of days, and a factor of ~10
variations in the solar wind flux can change the size of the
interaction region from a few to more than 20 RP.  It is therefore
critical to measure the solar wind for several solar rotations (~26
days per rotation) before and after the flyby in order to characterize
the most likely external solar wind properties during the actual
encounter period. Furthermore, since the strong asymmetry of the
interaction depends on the direction of the IMF, our analysis of SWAP
data will need assistance from increasingly capable models of solar
wind structure based on plasma and magnetometer data from spacecraft
elsewhere in the solar system.

2.1 Atmospheric Escape 

The exact nature of Pluto's plasma interaction is critically dependent
on the hydrodynamic escape rate of the atmosphere from its weak
gravity. Escaping neutrals are photoionized by solar UV (or, less
likely, suffer an ionizing collision). Freshly ionized particles
experience an electric field due to their motion relative to the IMF
(that is carried away from the Sun by the solar wind) and are
accelerated by this motional electric field, extracting momentum from
the solar wind flow. This electrodynamic interaction modifies the
solar wind flow. Estimates of Pluto's atmospheric escape rate, Q, vary
substantially: McNutt (1989) estimated 2.3-5.5 x 1027 s^-1 for
CH4-dominated outflow, Krasnopolsky (1999) found a hydrodynamic
outflow of N2 of 2.0-2.6 x 1027 s^-1, while Tian and Toon (2005)
derived values for N2 escape as high as 2 x 1028 s^-1. The nature of
the plasma interaction varies considerably over this range of escape
rates since the scale of the interaction is proportional to Q (Bagenal
and McNutt, 1989). 

Hydrodynamic escape of an atmosphere occurs when the atmospheric gases
in the upper atmosphere (in the vicinity of the exobase) are heated
significantly (so that the thermal speed is comparable to the local
sound speed). Krasnopolsky and Cruikshank (1999) reviewed the
photochemistry of Pluto's atmosphere while Krasnopolsky (1999)
summarized approaches taken in modeling the complexities of
hydrodynamic escape at Pluto. Earlier models approximated all the
heating that occurs in a thin layer of the atmosphere. Recently,
however, Tian and Toon (2005) developed a model that includes a
distributed heating function appropriate for EUV absorption by the
dominant molecule N2 relatively high in Pluto's atmosphere, which
leads to larger escape rates. These authors derive an exobase height
of 10-13 RP and transonic point of ~30 RP. The fact that even with
significant heating the escape speed (<100 m/s) is subsonic at the
exobase is consistent with what Krasnopolsky (1999) called "slow
hydrodynamic escape." Nevertheless, an exobase at 10-13 RP implies a
very extended atmosphere, and the present New Horizons trajectory with
a currently planned closest approach of ~9 RP, may well briefly dip
below the exobase. 

Tian and Toon (2005) modeled the effects of variations in (i) EUV flux
over the solar cycle, and (ii) Pluto's distance from the Sun, in order
to estimate the variability of atmospheric escape over the full Pluto
orbit. They found escape rates varying from 2 x 1028 s-1 (for Pluto at
30 AU and solar maximum activity) to 1 x 1028 s-1(for Pluto at 40 AU
and solar minimum). These authors, however, were not able to find a
stable solution for atmospheric escape at the low heating levels
appropriate for aphelion (50 AU). Whether Pluto retains a stable
atmosphere through aphelion or not remains an open issue. Recent
occultation measurements suggest that Pluto's atmosphere has been
expanding rather than collapsing as Pluto recedes from the Sun (Elliot
et al., 2003). The possibility that the atmospheric pressure could be
increasing on Pluto even after perihelion was predicted [Stern et al.,
1988; Hansen and Paige, 1996]. These authors suggested that there
would be a phase lag between perihelion and the maximum atmospheric
pressure. There is also the possibility that the polar obliquity is
now putting some new, fresh frost into sunlight, further driving
atmospheric expansion.

Assuming the current estimates for the escaping atmosphere have
persisted over the age of the solar system, Pluto could have lost
hundreds of meters to kilometers of material to space via "escape
erosion." In this case, the ancient topography made of volatiles (N2
and CO particularly, but also CH4) would have escaped to space. This
fantastic possibility appears to be something wholly unique to Pluto,
since even Triton is in the Jeans escape regime. Thus, escape erosion
might have caused Pluto to have a young surface, and if so, imaging
from New Horizons will give us information on the recent Kuiper Belt
impactor size-frequency distribution, in contrast to the time-integral
distribution provided by imaging of Charon, where such erosion is not
occurring.

2.2 Solar wind Interaction at High Atmospheric Escape

Galeev et al. (1985) derived a size for the interaction region (over
which significant momentum would be extracted from the solar wind) for
comets that is proportional to Q and inversely proportional to the
solar wind flux, nswVsw. Bagenal and McNutt (1989) applied this simple
scaling law to Pluto and found that for typical solar wind conditions
and photoionization of CH4 or N2, the scale size of the dayside
interaction region, RSO (the "stand-off" distance) is RSO/RP = Q/Q0)
where Q0 is 1.5 x 1027 s-1. Applying the upper range of atmospheric
escape rates discussed above, one finds RSO = 6-13 RP for 1-2 x 1027
s-1. (assuming an escape speed of 100 m/s). Note that these simple
calculations assume the exobase is close to the planet. Tian and Toon
(2005) suggest that the exobase may be as high as 10-13 Rp, in which
case the standoff distance would be expanded
accordingly. Nevertheless, the linear dependence on solar wind flux
means that larger (factor of ~10) variations in solar wind density
would produce a similar variation in RSO, often extending the
interaction region beyond the orbit of Charon at 17 RP.

The above simple scaling is based on a fluid approach that is
appropriate for cometary interactions in the inner solar system. The
weak IMF at 30 AU (see Table I) means that ion kinetic effects are
crucial at Pluto due to the large pickup-ion gyroradius (Kecskemety
and Cravens, 1993) and the large turning distance for solar wind
protons (Bagenal et al., 1997).

Previous models of weakly outgassing bodies used two-dimensional,
two-ion fluid and/or hybrid (i.e. fluid electrons, kinetic ions)
simulations (Bogdanov et al., 1996; Sauer et al., 1997; Hopcroft and
Chapman, 2001). All of these models showed the formation of asymmetric
plasma structures. Sauer et al. (1997) specifically investigated the
Pluto plasma interaction with two-dimensional models for Q > 1027
s-1. While the two-dimensional models provide a good qualitative
description of the interaction, the quantitative details of the plasma
coupling (i.e. momentum transfer) require three-dimensional
models. More recently, Lipatov et al. (2002) performed a
three-dimensional hybrid code simulation of the solar wind interaction
with weak comets and illustrated the dependence of gas production
rates (6.2 x 1026 s-1 < Q < 1028 s-1) on plasma structure. However,
these results are applicable to 1 AU where the ions are more strongly
magnetized compared to the situation at 30 AU where the IMF is very
weak and the solar wind is tenuous. 

Delamere and Bagenal (2004) modeled the kinetic interaction of the
solar wind with Pluto's escaping atmosphere using a hybrid simulation
that treats the pick-up ions and solar wind protons as particles and
the electrons as a massless fluid. A hybrid code is a reasonable
approach for a system with scale sizes ranging from the ion inertial
length of the solar wind protons (~2 RP) to the pick-up ion (N2+)
gyroradius (~500 RP). Figure 2 shows a schematic of the ion motion
near Pluto in this kinetic regime. Note the asymmetry of the
interaction imposed by the direction of the magnetic field in the
solar wind (Bsw). 

Figure 2. Schematic of ion motion near Pluto's interaction
region. Pickup ions move initially in the direction of the solar wind
convection electric field (+y) and the solar wind flow is deflected in
the -y direction, consistent with momentum conservation.

Figures 3, 4 and 5 show the density, temperature, and bulk flow
resulting from hybrid simulations for an atmospheric escape rate of Q
= 1 x 1027 s-1 |B| = 0.2 nT and an atmospheric escape speed of 100 m/s
(Delamere and Bagenal, 2004) in the same geometry depicted in Figure 2
(upstream solar wind flow from the left and IMF pointing out of the
page). The main features of the interaction are a broad region of
proton heating extending ~20 RP upstream of Pluto consistent with the
expected location of a bow shock. The kinetic energy of the solar wind
protons is not fully converted to thermal energy, and the shock
structure forms only in the upper half of the interaction region. In
the lower half of the interaction region the protons are compressed
(illustrated by higher density in Figure 3, higher temperature in
Figure 4, slightly reduced and deflected flow in Figure 5) behind what
is effectively a magnetic pile-up boundary. Since the gyroradii of the
pick-up ions is huge compared with the scale of the interaction
region, they effectively move upwards in the geometry illustrated in
these figures. The density of pick-up ions is relatively small but,
because of their high mass, the momentum they extract from the solar
wind is significant. The solar wind is slowed down (leading to
compression and enhanced magnetic field) and deflected in the opposite
direction to the pick-up ions. Superimposed on the density contours in
Figure 5 are trajectories of solar wind protons. The large gyrations
reflect the weak magnetic field's inability to turn the protons. The
smaller gyrations indicate stronger magnetic fields in the pile-up
region. Harnett et al. (2005) showed that multi-fluid simulations
produced qualitatively similar, though less pronounced, asymmetries in
the interaction region.

Figure 3. Sample trajectories of solar wind protons near Pluto (origin
of plot) through the interaction region with total ion density (N2+
plus solar wind proton).  The trajectories are altered at the magnetic
pileup boundary, but only partially thermalized as illustrated in
Figure 4.  The magnetic pileup boundary is very abrupt in the -z half
plane.

Figure 4. Column averaged proton temperature in the plane
perpendicular to Bsw.   The maximum temperature for the solar wind
protons (i.e. fully thermalized) moving initially at 340 km/s is 360
eV.  The interaction only partially thermalizes the solar wind with
maximum temperatures of a few 10s of eV.  The dashed line is the
expected location (i.e. based on upstream standoff distance) of the
bow shock and the dash-dot-dot-dot line is the expected location of
the magnetic pileup boundary based on calculations for solar wind
stagnation near comets for Q = 1027 s-1 and Bsw = 0.2 nT (Biermann
(1967), Galeev et al.,  (1985)).   Given the extreme asymmetry of the
interaction, it is not possible to fit the bow shock and the magnetic
pileup boundary with symmetric parabolic functions.  We based our fit
for the bow shock on the +z half plane and the fit for the magnetic
pileup boundary on the -z half plane.

Figure 5. Ion bulk flow velocity (pickup ions plus solar wind protons)
with contours of total ion density illustrating the high degree of
asymmetry in the plasma flow.

Clearly, the upstream direction and magnitude of the IMF are critical
drivers for the solar wind interaction with Pluto. Unfortunately, New
Horizons did not have the resources to carry a magnetometer, so we
will need to use our best knowledge of solar wind physics to infer the
magnetic properties. Generally speaking, the IMF becomes increasingly
tightly wound spirals on conical surfaces (for different latitudes)
with increasing distance from the Sun. Thus, we expect the average
field direction to be nearly perpendicular to the flow by 32
AU. Because of the alternating sector structure of the IMF, it is
equally likely that the IMF will be pointed in either
direction. Measurements of pick-up ions, either by SWAP or by PEPSSI
[McNutt et al., 2007] will provide the critical indication of the sign
of the magnetic field at the time of the encounter.

2.3 Solar wind Interaction at Low Atmospheric Escape 

If Pluto's escape rate is less than ~1027 s^-1, then the solar wind is
expected to penetrate close to Pluto and impinge directly onto its
ionosphere. Strictly speaking, the "obstacle" that deflects the solar
wind in this interaction is produced by the electrical currents
induced in the ionosphere. Such an interaction would be similar to the
solar wind interaction with Venus or the Saturnian magnetospheric
interaction with Titan (Luhmann et al. 1991). The photochemistry of
Pluto's upper atmosphere and ionosphere modeled by both Krasnopolsky
and Cruikshank (1999) and by Ip et al. (2000) show peak ionospheric
densities of a few x 103 cm-3 at altitudes of ~1000 km above the
surface of Pluto. They find the main ionospheric ion to be H2CN+. 

The solar wind deflection and any pick-up ions will be significantly
harder to measure on the New Horizons flyby in the case of a
Venus-like interaction. However, the broad nature of the bow shock and
the large gyroradii of any ions beyond the ionosphere should still
produce detectable signatures several radii away from Pluto.

2.4 Solar wind Interaction at Aphelion

It is not clear whether the atmosphere of Pluto completely collapses
onto the surface when Pluto reaches aphelion. In the absence of a
significant atmosphere Pluto will have a "Moon-like" interaction, as
is expected for its moon Charon. In this type of interaction, the
solar wind is absorbed on the sunlit side and the IMF diffuses through
the non-conducting bodies, generating an extremely hard vacuum in the
cavity behind.  It seems unlikely that at temperatures below 40K
Pluto's icy outer layers could be electrically conducting. But should
they have significant conductivity, the plasma interaction may be
similar to the solar wind/asteroid interaction described by Wang and
Kivelson (1996), Omidi et al. (2002), and Blanco-Cano et
al. (2003). If the dimensions of the obstacle are small compared to
the ion inertial length, then the interaction is mediated by the
whistler mode. Comparison of Galileo observations of
asteroid-associated perturbations with numerical models confirm the
whistler-mode interaction. Magnetohydrodynamic shock waves are absent
with small obstacles as compressional waves can only exist on size
scales larger than the ion inertial length. Pluto represents a
possible intermediate case if the interaction region is limited to an
area close to the planet as the solar wind proton inertial length is
roughly 2 RP. 

2.5 Heliospheric Pickup Protons

In addition to the primary science of measuring the solar wind
interaction with Pluto, SWAP may afford an excellent opportunity to
measure heliospheric pickup protons on its way out through the
heliosphere, en route to Pluto (and beyond).

The solar system moves continually through the local interstellar
medium (LISM) - the part of the galaxy nearest to our solar system -
causing a complex set of interactions between the outflowing solar
wind and the matter in the LISM (see Figure 6). These interstellar
interactions set up the plasma boundaries of our heliosphere: the
termination shock where the solar wind abruptly slows, becoming
subsonic, prior to bending back into an extended tail in the
interstellar downstream direction; the heliopause separating the
subsonic solar wind from the ionized plasma of the interstellar
medium; and possibly a bow shock where the LISM plasma becomes
subsonic prior to being deflected around the heliosphere. 

Figure 6. The SWAP instrument will measure protons produced from
neutral matter that drifts into the heliosphere from the LISM.

Interstellar atoms, predominantly H, continually drift into and
through the heliosphere, and due to their neutrality are unimpeded by
the solar wind. The H atoms on trajectories toward the Sun move into
regions of increasingly dense solar wind and higher levels of solar
radiation, enhancing the probability of ionization through
charge-exchange or photo-ionization. The vast majority of interstellar
H atoms that penetrate inside of 4 AU become ionized and incorporated
into the solar wind flow. Upon ionization, newborn ions begin
essentially at rest in the reference frame of the spacecraft and
Sun. Just like for planetary pickup ions, as described above, because
the moving solar wind carries frozen-in magnetic field lines, the
newborn ions encounter a motional electric field (-vsw x B) exerted by
the solar wind and become picked-up and carried outward with the solar
wind. As they move outward, the pickup ions also scatter due to
magnetic inhomogeneities in the IMF, forming an almost spherical ring
distribution in velocity space. As measured by spacecraft, these
interstellar pickup-ion distributions are essentially flat for speeds
less than 2 vsw, then drop off sharply at speeds above this limit
(Gloeckler et al., 1995).

The first interstellar pickup ions discovered (Mobius et al., 1985)
were He+ created by ionization of interstellar neutral He that
penetrated within 1 AU of the Sun. The composition and velocity
space-resolved measurements by the SWICS experiment on Ulysses
(Gloeckler et al., 1992) made it possible to explore pickup ions from
1.35-5.4 AU in great detail. The review by Gloeckler and Geiss (1998)
provides an excellent summary of these pickup-ion observations. The
observations include the most abundant pickup ion, H+, second most
abundant, He+, and several other interstellar pickup-ion species, N+,
O+, and Ne+ (Geiss et al., 1994). The SWICS team demonstrated the
existence of pickup distributions of He++, which is produced largely
by double charge exchange of atomic He with solar wind alpha particles
(Gloeckler et al., 1997) and rare 3He+ pickup ions. In addition to the
interstellar pickup ions, SWICS distributions showed that the majority
of the C+ and a fraction of the O+ and N+ are produced by an
additional "inner source" of neutral atoms located near the Sun (Geiss
et al., 1995). The inner-source, velocity-space distributions are
significantly modified as they cool over the solar wind's transition
from the near-Sun source to several AU where they were observed
(Schwadron et al., 2000). 

Ulysses/SWICS also observed the ubiquity of pickup-ion "tails" in slow
solar wind (Gloeckler et al., 1994; Schwadron et al., 1996; Gloeckler,
1999). These tails do not correlate strongly with the presence of
shocks but with compressive magnetosonic waves, suggesting that pickup
ions are subject to strong statistical acceleration through processes
such as transit-time damping of magnetosonic waves in slow solar wind
(Schwadron et al., 1996; Fisk et al., 2000). 

Prior to Ulysses, it was expected that pickup-ion distributions should
be fairly isotropic due to pitch-angle scattering from background
turbulence and self-generated waves (Lee and Ip, 1987). Instead,
pickup-ion distributions were observed to be highly anisotropic
(Gloeckler et al., 1995) with scattering mean free paths ~1 AU; the
most likely cause is the inhibition of scattering through 90 degrees pitch
angle (Fisk et al., 1997). Although this lack of scattering is not
fully understood, it has been shown that the turbulence of the wind
has a strong 2-D component (Matthaeus et al., 1990; Bieber et al.,
1996) that is ineffective for pickup-ion scattering (Bieber et al.,
1994; Zank et al., 1998).  Pickup-ion models have been devised that
take into account the long-scattering mean free path (e.g., Isenberg,
1997; Schwadron, 1998). 

Interestingly, the same self-generated turbulence associated with
pickup-ion scattering has been considered an energy source for heating
the solar wind (Zank et al., 1996; Matthaeus et al., 1999). These
theories for the turbulent heating of solar wind take into account
both the pickup-ion driven turbulence, which is predominant outside of
~8 AU, and wind shear. Smith et al. (2001) and Smith et al. (2006)
have rigorously tested the heating models using data from Voyager 2
and Pioneer 11, and improved models of the pickup-ion scattering and
associated wave excitation (Isenberg et al., 2005) have lead to
remarkable agreement between data and theory.  

Charge-exchange between interstellar hydrogen and solar wind protons
leads to a complex interaction near the nose of the heliosheath, where
a so-called hydrogen wall is formed from slowed interstellar hydrogen
atoms and charge-exchanged solar wind protons. These interactions
cause the removal or filtration of a fraction of the penetrating
interstellar hydrogen atoms (e.g., Baranov and Malama, 1995). The
solar radiation pressure and the rates of photo-ionization and charge
exchange vary with solar latitude and over the solar
cycle. Sophisticated models of interstellar neutral atoms have been
developed to take these effects into account (e.g., Izmodenov et al.,
1999). 

Cassini was the first spacecraft to measure pickup ions in the
interstellar downstream region of the heliosphere beyond 1 AU (McComas
et al., 2004a). Both interstellar pickup H+ and He+ were identified
with the familiar cutoff in velocity space at twice the solar wind
speed. Observed enhancements in the pickup He+ were consistent with
gravitational focusing by the Sun. Further, McComas et al. also
reported observations of the interstellar hydrogen shadow caused by
depletion of H atoms in the downstream region due to the outward force
of radiation pressure (which exceeded the gravitational force at the
time of observation) and the high probability of ionization for atoms
that must pass close to the Sun to move behind it.  

The trajectory of New Horizons takes the spacecraft out to Pluto,
which is currently toward the nose of the heliosphere. Figure 7 shows
simulated distribution functions of pickup protons and solar wind
protons in the spacecraft frame. We have used a kappa distribution for
the solar wind protons (e.g., Vasyliunas, 1968; Collier, 1993) (with a
kappa ~ 3, which is typical), and a density that falls off as R-2 from
a 1 AU density of 5 cm-3 and temperatures consistent with observations
(e.g., Smith et al., 2001). For the pickup protons, we have taken a
steady-state distribution that includes convection and adiabatic
cooling in the radially expanding flow (Vasyliunas and Siscoe,
1976). The proton-pickup rate at each location is the local ionization
rate times the neutral density solved using the "hot" model (Fahr,
1971; Thomas, 1978; Wu and Judge, 1979), which accounts for
gravitational focusing by the Sun, ionization loss, and the finite
temperature of incoming neutral atoms. We have taken an interstellar
density near the termination shock of 0.1 cm-3, and neutral
temperature of 11 000 K, a neutral inflow speed of 22 km/s, an
ionization rate of 7 x 10-7 s-1 (referenced at 1 AU). We have also
taken the force of radiation pressure to be comparable to that of
gravity. The top panel in Figure 7 shows that solar wind protons are
typically comparable to pickup protons near 5 AU. The pickup protons
are observable near the knee of the distribution, so long as the solar
wind is not significantly hotter or has significantly broader tails
than those shown. Since the plasma near the ecliptic plane is highly
variable near 5 AU, the detection of pickup protons is possible, but
often obscured by the solar wind proton distribution. As New Horizons
moves out toward the nose of the heliosphere, the situation
changes. The solar wind density falls off with the square of
heliocentric distance (while the pickup proton density falls off only
as R-1) and cools. Both of these effects make it much easier to detect
"clean" pickup-proton distributions.

Figure 7. Pickup-proton (solid curves) and solar wind (dashed curves)
distribution functions shown with distance from the Sun toward the
nose of the heliosphere. The pickup-proton distribution functions
become more and more prominent compared to the cooled solar wind
distributions further out in the heliosphere.    

The unique trajectory of New Horizons combined with SWAP's capability
of pickup-proton measurement should enable it to address some
fascinating issues. The distribution functions measured by SWAP will
allow investigation of possible sources (other than interstellar) of
pickup ions. For example, it is thought that an Outer Source of pickup
ions may be caused by the interaction of solar wind with dust grains
from the Kuiper Belt (Schwadron et al., 2002). The Outer Source is
likely to be far more variable in space and time than the interstellar
source.

3.0 Instrument Description

The SWAP electro-optics consists of 1) a retarding potential analyzer
(RPA); 2) a deflector (DFL); and 3) an electrostatic analyzer
(ESA). Collectively, these elements select the angles and energies of
the solar wind and pickup ions to be measured.  Ions selected by the
electro-optics are then registered with a coincidence detector
system. Figure 8 schematically depicts the SWAP principle
electro-optics. Ions enter through the RPA with all ions having energy
per charge (E/q) less than the RPA voltage being rejected by this
hi-pass filter. Ions entering at angles from above horizontal in this
figure can be deflected in to the subsequent electro-optics by
applying a voltage to the deflector ring. Ions with E/q greater than
the RPA voltage are then selected by the ESA, which rejects ions
outside the selected E/q range as well as UV light and neutral
particles.

Figure 8. Schematic diagram of the SWAP electro-optics including RPA,
deflector, ESA, and detector section.

Figure 9 schematically shows how SWAP's ESA and RPA are used together
to select the E/q passband. When the RPA is off, the passband is
determined solely by the ESA, which has an 8.5% FWHM resolution (top
panel). At increasing RPA voltages for a given ESA setting, the
passband is cutoff in a variable "shark-fin" shape, allowing roughly
two decades decreased sensitivity (middle panel). Finally,
differentiating adjacent RPA/ESA combinations, or better yet
deconvolving multiple combinations, provides high resolution
differential measurements of the incident beam.

Figure 9. SWAP principle of operation. Different combinations of RPA
and ESA settings provide for a variable E/q passband. Multiple
combinations of settings can be differentiated to produce very
high-resolution measurements.

As an example of how the RPA can be used to measure very small changes
in beam energy, and hence solar wind speed, Figure 10 shows the count
rate as a function of RPA voltages for a range of different ion beam
energies taken with the flight electro-optics during instrument
calibration (described in more detail in Section 4.1).  A small subset
of the full energy range of SWAP (990 to 1010 eV) is shown to
highlight the energy resolution possible by this instrument. For each
scan, the count rate is normalized to the rate when the RPA voltage is
set at zero to take into account the differences in the ion beam
flux. Differences in the ion energy as small as 1-2 eV are
distinguishable at typical solar wind energies of ~1000 eV.

Figure 10. RPA resolution of ion beams with energies from 990-1010
eV. Each RPA scan shown was taken with the beam incident normal to the
RPA. Deconvolution of such curves should make changes in beam energy
as small as 1-2 eV possible.

The transmitted ions are post-accelerated into the detector section,
which employs an ultra-thin carbon foil and two channel electron
multipliers (CEMs) to make a coincidence measurement of both the
primary particle and the secondary electrons generated when the
primary particle passes through an ultra-thin carbon foil. Charge
amplifiers (CHAMPs) service the two CEMs and transmit digital pulses
when events are detected.  High voltage power supplies (HVPS) provide
power for the CEMs and sweep the voltages on the electro-optics.  The
control board processes the pulses from the CHAMPs, controls the
sweeping of the high voltages, digitizes the housekeeping data,
creates telemetry packets for transmission, accepts commands, and
converts spacecraft power into the secondary voltages required by the
instrument. Key properties of the SWAP instrument are given in Table
II. 

Table II. Measured properties of the SWAP instrument

The New Horizons spacecraft points its optical instruments by rotating
the spacecraft about its Z-axis. In order to make solar wind
measurements largely independent of this scan angle, SWAP was designed
to be essentially symmetric about Z. SWAP is mounted into a
cylindrical bracket as depicted in Figure 11. By placing it on the -Z
corner of the spacecraft, the instrument's field-of-view is clear of
any structures or obstructions for all scan angles.

Figure 11. Location of SWAP on the New Horizons Spacecraft 

3.1 Electro-Optic Design

The electro-optic design is driven by the unique needs of the New
Horizons Mission to Pluto.  A large aperture instrument is required
since the density of the solar wind falls off as the square of the
distance from the Sun.  At ~32 AU, where we will encounter Pluto, the
mean solar wind density is roughly 3 decades lower than the solar wind
at 1 AU.  Also, the New Horizons spacecraft does not have a scan
platform for the optical instruments.  Instead, the spacecraft slews
about the Z-axis to provide the required scanning motion.  Since
SWAP's focus is to measure the solar wind, we required an
electro-optic design that was largely unaffected by this rotation
about the Z-axis.  Therefore, we selected a top-hat design with an ESA
that has a large angular and energy acceptance to maximize the
effective aperture.  Since the driving measurement is to look for
small changes in the solar wind speed as New Horizons passes Pluto, we
included an RPA to allow finer resolution energy measurements and to
provide a variable electrostatic passband. This is accomplished by
"crossing" the ESA and RPA such that the RPA admits particles only in
the highest energy portion of the ESA passband as described above.
Also, because the spacecraft will sometimes be required to tip out of
the ecliptic plane, we incorporated a deflector to adjust the look
direction of the instrument.  Finally, because the mission is designed
so the spacecraft rarely points the - Y-axis toward the Sun, we can
provide an unobstructed field of view of >270 degrees in the X-Y plane,
centered on the +Y-axis, which is aligned with the high gain antenna
(HGA).  Nominally, the HGA points toward the Earth within a few
degrees of the Sun for a large fraction of the Pluto encounter. All
supports for the top-hat top-plate, RPA, ESA, and the aperture doors
for protection during launch are restricted to the remaining
90 degrees aligned with the - Y-axis.  Figure 12 shows the key dimensions of
the SWAP electro-optics. 

Figure 12. Key dimensions of the Electro-Optics

3.1.1 RPA Design

The RPA consists of four concentric aluminum cylinders or screens,
each of which is machined with ~90,000 close-packed holes to create a
self-supporting grid structure.  Each cylinder, approximately 65%
transmissive, is 0.762-mm thick and has been machined into a grid-like
structure by drilling 0.343-mm-diameter holes through nominal
0.394-mm-thick aluminum in a close-packed hexagonal configuration
(Figure 13). From outermost to innermost, the outer diameters of the
four RPA screens are 174.4, 169.6, 166.4, and 161.6 mm. The outside
and inside cylinders are at ground potential (0 V). The two central
cylinders are biased from 0 to +2000 V in 0.49 V steps and isolated
from the rest of the structure by ceramic insulators.  There is a
2.032 mm gap between the biased and grounded RPA grids to provide
sufficient high voltage clearance for the -2000 V potential.  The two
biased RPA grids are separated by 1.016 mm to stop penetration of
lower energy equipotentials through the grid holes, thus creating a
much more uniform electric field in the RPA region. The unobstructed
Field of View (FOV) is 276 degrees in the roll direction.

Figure 13. Photograph of 0.343 mm RPA holes on 0.381 centers

The RPA provides a low-pass filter with a relatively sharp energy
cutoff so that we can make a fine sweep across the solar wind beam
after locating it with a coarse ESA energy scan.  Ions that have
sufficient energy to climb the electrostatic "hill" set by the voltage
on the inner RPA grids are reaccelerated to their original energy as
they pass from the inner RPA grids to the final grounded RPA grid.

3.1.2 Deflector

SWAP incorporates a very simple deflector that is used to deflect
particles from above the central plane of the instrument (from further
out in the -Z axis of the spacecraft) into the ESA.  The deflector is
located just inboard of the RPA.  The voltage on the deflector ring is
varied from 0 to +4000 V.  It deflects up to 7000 eV/q particles up to
15 degrees into the ESA (larger deflections for lower energies).

3.1.3 ESA

The ESA provides coarse energy selection and protects the detectors
from UV light.  The dimensions of the top-hat ESA are shown in Figure
12.  The outer ESA is serrated and blackened with Ebanol-C, a
copper-black process that greatly reduces scattering of light and
particles.  The inner ESA, which is blackened but not serrated, is
supported on insulators that attach it to a cantilevered support
structure (Figure 14).  On the other side of this structure a grounded
cone completes the ESA design by providing a field-free region for
particles to enter into the detector region.  The voltage on the ESA
is varied from 0 to -4000 V.  It has a ratio of central E/q of
particle transmitted to ESA voltage, or "K-factor", of 1.88 and
selects up to 7.5 keV/q particles (passband central energy) with a
delta (E)/E of 8.5% FWHM.

Figure 14. Cantilevered Inner-ESA dome with (bottom panel) and without
(top panel) blackening.

3.2 Detector Design

SWAP employs a coincidence system to detect incoming ions.  After ions
of the desired energy and angle have been selected by the
electro-optics system, they pass through a field free region between
the ESA and detector region.  Once a particle enters the detector
region, it is accelerated by the CEM high voltage towards the Focus
Ring on which is suspended an ultra-thin carbon foil (McComas et al.,
2004b).  The carbon foil is nominally 1 micro g/cm2 thick and is suspended
on a 64% transmissive grid.  The particle passes through the Focus
Ring, which is at the Primary CEM (PCEM) HVPS output voltage, and
travels on to the PCEM (Figure 15).  Forward-scattered electrons from
the carbon foil are also accelerated to the PCEM due to the ~100 V
potential created by the PCEM strip current and a resistive divider.
Backward scattered electrons are directed by the Focus Ring towards
the Secondary CEM (SCEM) which collects them.

Figure 15. PCEM and Focus Ring with Carbon Foil, partially assembled.

Counts from these two CEMs are registered by CHAMPs and their
associated electronics.  A count from either the PCEM or the SCEM
starts a 100 ns coincidence window timer.  Given the particle
trajectories and the electron trajectories, it is possible for either
the PCEM or the SCEM to trigger first, so no specific arrival order is
required by the electronics. The CEMs (Figure 16) have quad spiral
channels with a resistance of 300 Mohm and dark counts less than 0.04
s^-1.

Figure 16. Quad Spiral Channel Electron Multiplier

The coincidence detector system reduces the background from CEM dark
counts, penetrating energetic particles and UV noise, allowing SWAP to
have a very low enough noise floor.  Two detectors further provide the
redundancy needed for a long-duration mission.  SWAP can still make
its primary science measurement using either one of its CEMs.  This
requirement for redundancy drove us to add a -1 kV Focus Ring supply
that is slaved to the SCEM and diode OR-ed with the PCEM supply to the
Focus Ring.  If the PCEM is turned off, the -1 kV Focus Ring supply
will accelerate back-scattered electrons from the carbon foil into the
SCEM.

The long CEM lifetime needed for this mission required us to select
all materials near the detector to be ultra-low outgassing.  We used
only glass, metal, and ceramic materials for all parts of the detector
that had venting access to the CEM detectors.  HV and LV cabling was
brought to bulkheads, and the signals were conducted through ceramic
feedthroughs to the exact location they were required (Figure 17).

Figure 17. Detector Wiring Diagram

3.3 Mechanical Design

The SWAP mechanical design consists of three main subassemblies:
optics and detector, aperture door, and electronics packaging and
cabling (Figure 18). In addition to these instrument subassemblies, we
also discuss SWAP's structural and thermal design in the following
sections.

Figure 18. Photograph of SWAP instrument highlighting subassemblies

3.3.1 Optics and detector mechanical design

The critical requirements driving the optics and detector mechanical
design were to mount the electro-optical components as specified by
the ray-traced model, produce optically black surfaces inside the
instrument, ensure contamination control for the detectors, provide
for easy refurbishment, and accurately align all critical components.

We used a SIMION ray tracing model to define the electro-optics
component (Figure 19) sizes and locations (RPA, ESA, CEMs), and
required that the mounting of the various elements should minimally
impact particle trajectories.  We accommodated this by mounting the
Inner ESA and Secondary CEM assemblies from a cantilevered support
hidden in the 90 degrees region where SWAP does not require particle viewing.
The location of this cantilever matches up with the hinge assembly on
the door.  

Figure 19. Mechanical configuration of the optics and detector
assembly.

In order to accommodate the cleanliness requirement, the optics and
detector areas included a split between ultra-clean and electronics
volumes.  We used ceramic for the electronics near the primary CEM,
and fed the cabling into the primary CEM area using pass-thrus in the
cantilever support so that it would not pass through the ESA gap.

We refurbished SWAP after spacecraft environmental testing, replacing
the CEMs and the ultra-thin carbon foil with ones that had been stored
in a clean environment after burn-in.  To minimize the risk of
dismantling the instrument, the CEM assemblies were designed so that
they could be easily replaced. In order to meet the alignment
requirements for the electro-optics, we designed features that would
control the concentricity and placement of the optics and detectors in
relation to each other both before and after refurbishment.

3.3.2 Aperture door design

The aperture door's (Figure 20) main requirement is to protect the
SWAP RPA from contamination and damage during ground and launch
operations.  Although we designed the door to open one time after
launch, it was important that it easily reset for test on the
spacecraft.  We opened the door multiple times as part of the
instrument verification and during the environmental test flow at the
spacecraft level.  The door could not be a glinting source because it
stays with SWAP during the entire mission to Pluto. Spacecraft
requirements dictated that external surfaces be conductively coupled
to the instrument. Therefore, we coated the door surfaces with black
nickel and provided for ground straps built into the design.  The New
Horizons Component Environmental Specification required verifying
through test that the door's torque margin was greater than 2.25.
Testing demonstrated a torque margin of 3.10, and the door
successfully opened in flight.

Figure 20. Door assembled on the SWAP outer ESA. The TiNi Pinpuller
releases the door so it can spring open.

3.3.3 Electronics packaging and cabling design

Figure 21 shows the as-built electronics packaging and cabling that
are part of the SWAP instrument.  Panel a) shows the HVPS and Control
board with the electronics before the electronics housing was
installed.  Panels b) and c) show different views of the CHAMP
electronics and cables that bring the HV to the RPA and ESA.

Figure 21. As-built electronics packaging and cabling.

To save weight and volume, the electronics are compactly packaged
below the ESA in an "electronics volume", as shown in panel a).  The
electronics are laid out on three boards, which are connected via
pin/socket connections.  The top board contains the HV, the middle
board includes the drivers for the HVPS, and the bottom board receives
power from and provides telemetry to the spacecraft.  The CHAMP boards
were originally intended for the lower electronics, but engineering
testing showed that they needed to be closer to the detectors to
decrease noise.  We routed the HV and signal cables through the
electronics housing to their appropriate location and included cable
routing tie points in the outer ESA design to prevent the cables from
moving during launch loads or transportation.  The standoffs between
the three bottom boards structurally attached the boards to the
housing and allowed heat to conductively pass to the housing.

3.3.4 Structural Design

The Component Environmental Specification defined the following
critical structural requirements for SWAP: 1) quasi-static load of 30
gs force applied separately along 3 orthogonal axes, 2) first mode
structural frequency constraints of  >70 Hz (thrust direction) and
(>50 Hz) (lateral direction),  3) sine vibration up to 20 g (22-24 Hz,
and 4) Random vibration with an overall amplitude of 10.4 Grms.

We performed a structural analysis during the final design phase to
very that the SWAP structure would meet these requirements.  The
analysis focused on the critical items such as the cantilevered ESA
and outer ESA. The analysis also showed that the first natural
frequency of the SWAP instrument, 180 Hz, was well above the
requirement.  Finally, we performed vibration testing during
environmental testing to ensure that SWAP would survive the prescribed
levels.

3.3.5 Thermal design

SWAP has flight temperature limits of 0 to +40 degrees C (operating) and -20
to +50 degrees C (non-operating). We initially performed a thermal analysis to
show that the electronics and other temperature-sensitive parts of the
SWAP instrument could survive these temperature extremes.  The
analysis also showed that the heat exchange with the spacecraft was
less than 5 Watts (the limiting case was when the instrument was at
its hot limit with the door open), as required by the spacecraft.
During environmental testing, we performed a thermal vacuum test to
validate the functionality of SWAP at all temperatures within these
extremes; hot and cold turn-ons of the instrument were also performed.

3.4 Electronics

3.4.1 CHAMPs

The CHAMPs for SWAP convert a charge pulse from the CEMs into a logic
pulse that can be registered and processed on the Control Board. The
CHAMPs (Figure 22) are located as close to the detectors as practical
in separate enclosures mounted to the top of the strong back.  The
SWAP design incorporates high-speed commercial hybrid CHAMPs, which
work reliably to >1 MHz rates and have voltage-adjustable
thresholds. The output from the CEMs is brought to the CHAMP through a
short coaxial cable through a bulkhead connector that ties the shield
to chassis ground.  A 100-ohm resistor and back-to-back diodes protect
the input from discharges, and the input is AC coupled to the charge
amp.  

Injecting test pulses into the front end demonstrates the end-to-end
integrity of the microcircuits and cabling for CEM pulse processing.
Test pulses are driven from the Control Board to the CHAMP, which is
twisted with a return line.

The threshold voltage for the CHAMP is set by a Digital-to-Analog
Converter (DAC) on the Control Board, and this value can be set
through standard software commands, allowing the threshold to be
updated at any time.  A resistor sets the output pulse width to 70 ns
and the amplifier dead-time to 100 ns.  The output pulses from the
amplifier are buffered by two Schmitt trigger buffers and transmitted
through a back-terminated series resistor to the cable that connects
to the Control Board.

Figure 22. Charge Amplifier Photograph

3.4.2 HVPS

The HVPS (Figure 23) set the voltages on the optical surfaces (RPA,
DFL, ESA, & Focus Ring) as well as supplying power to the Primary and
Secondary Channel Electron Multipliers (PCEM and SCEM). Table III
shows the primary HVPS properties. The two detector supplies are
single string, but the two independent detectors provide redundancy.
The Focus Ring is diode-ORed with the output of the PCEM and a -1000 V
Focus Ring supply that is created and controlled in parallel with the
SCEM.  The optical power supplies are redundant and diode-ORed
together. The secondary ground on the HVPS is "zap-trapped" to chassis
through back-to-back diodes for protection. Mechanically, the HVPS are
built onto two interconnected boards: the Driver Board and the
Multiplier Board.

Figure 23. HVPS being assembled into the SWAP sensor

Table III. Key HVPS Specifications.

3.4.3 Control Board

The SWAP Control Board provides the electrical interface between the
New Horizons spacecraft and the SWAP instrument.  All command,
telemetry, power, safe/arm, and actuation interfaces reside on the
Control Board.  Software on an 8051 microcontroller responds to
commands, controls the operation of the instrument, sequences the high
voltage power supplies, collects the data, and formats telemetry for
down-link.

SWAP communicates to the spacecraft through two redundant (A & B side)
asynchronous RS-422 interfaces.  SWAP accepts serial commands,
produces serial telemetry, and synchronizes communication with the
spacecraft through a Pulse Per Second (PPS) line.  The command and
telemetry data are transmitted at 38,400 Baud.

The Control Board contains the EMI filters and DC-DC converters, which
create the +5 and -5 Volt secondary power rails that are isolated from
the spacecraft primary bus.  Two separate 1.5W DC-DC Converter provide
power for the instrument.  Since the DC-DC converters do not have
provisions for synchronization, each converter has an independent EMI
filter to eliminate low-frequency noise from the beat between the two
converter oscillators.  The Control Board has a set of power MOSFETs,
which allow us to switch low voltage power to the PCEM HVPS, SCEM
HVPS, Optical HVPS Bank A, Optical HVPS Bank B, and the Housekeeping
circuits.

The micro-controller provides all of the on-board processing required
to operate the instrument, responds to commands, and produces
telemetry.  The CPU executes at 4.9152 MHz, and a minimum of 12 clocks
are required to execute a single instruction, so the top speed of the
processor is 0.4 million instructions per second (MIPS).  From this
low rate, the clock can be divided down further to reduce the
instruction rate to 0.05 MIPS.  Due to spacecraft-level power
constraints every effort was made to reduce the operating power
required for SWAP.  This maximized SWAP's ability to remain on during
the Pluto encounter sequence when the spacecraft is power-limited.

Boot code for the micro-controller resides in a radiation hardened 32k
x 8-bit Programmable Read-Only Memory (PROM).  This ensures that the
instrument can always boot and establish basic communication with the
spacecraft even if other memory devices have endured temporary upsets
or even permanent degradation.  Two separate 128k x 8-bit Electrically
Erase-able Programmable Read-Only Memory (EEPROM) devices provide
redundant storage for a 64k x 8-bit storage area for program code an a
64k x 8-bit storage area for Look-Up Tables (LUTs).  Finally, a 128k x
8-bit Static Random Access Memory (SRAM) is used to provide 64k x 8
bits of code memory and 64k x 8 bits of data memory.  During normal
operations, if the program code in one of the Electrically Erasable
Programmable Read-Only Memory (EEPROM) banks has a valid checksum,
then it is loaded into RAM by the Boot PROM and executed.  A
Field-Programmable Gate Array (FPGA) controls the memory mapping and
memory windows that the micro-controller needs to access these
devices.

Signals received from the primary and secondary CHAMPs are processed
on the Control Board.  When the electro-optics have been set to the
appropriate levels and sufficient settling time has elapsed, the
software opens up an acquisition window and totals all of the primary
and secondary CEM pulses that occurred during the acquisition window.
Whenever a secondary or primary event occurs, a 100 ns coincidence
window is started.  If a pulse is received from the other charge
amplifier during the 100 ns coincidence window, then a coincidence
counter is incremented.

The Control Board sets the RPA, DFL, and ESA high voltage optical
settings using a set of independent digital-to-analog converters
(DAC).  The analog settings, along with digital enables, and the low
voltage power for the primary and redundant supplies are carried to
the HVPS over a dedicated board-to-board connector.  The HVPS returns
the analog current and voltage monitors to the Control Board for
analog housekeeping and real-time monitoring.

The health and safety of the instrument in general and the CEMs in
particular are monitored extensively by the Control Board.  The count
rate of each detector is totaled and checked against a software limit
every 0.5 s.  The supply voltage and strip current for each CEM is
monitored and verified by the software every 1 s.  The strip current
is also compared against a threshold level with 32 values.  If the
strip current exceeds this threshold, the software receives an
interrupt and may immediately take action.  This information, along
with temperatures in the instrument and the current and voltage
monitors of the low voltage power supplies are transmitted to the
ground in housekeeping telemetry packets.

3.5 Modes of Operation

SWAP flight activities are conducted during the following operational
phases: 1) Commissioning; 2) Jupiter Encounter; 3) Cruise/Annual
Checkout; and 4) Pluto Encounter.  To accommodate these different
activities, SWAP has various modes that are controlled by the SWAP
onboard flight software (FSW).  A mapping of the operational phases to
SWAP modes is described in Table IV, and a description of the modes
and corresponding power and telemetry resource consumption is shown in
Table V.

Table IV. Operational Phases and SWAP Modes

Table V.  SWAP Modes and Resources

Mode changes are triggered by command with the exception of the BOOT
to Low Voltage Engineering Mode transition, which is based on time
since power up, and is automatic.  Flight operations, and therefore
the modes, drive the requirements for the FSW.

3.6  Flight Software

3.6.1 Overall Capabilities

The FSW runs on the 80C51 microprocessor at 4.9152 MHz, equating to
approximately 400,000 instructions per second.  It is written almost
entirely in the C programming language.  The FSW consists of two
different code images stored in nonvolatile memory. At power on, the
FSW runs directly from a programmable, read-only memory (PROM) code
image.  This code gathers diagnostic information on the SWAP hardware,
provides the ability to reload EEPROM code or tables and determines
which copies of EEPROM code or table to use for science acquisition.
Once the EEPROM code and tables have been chosen, the EEPROM code is
copied into RAM for execution.

The FSW addresses these requirements areas: data interfacing and
synchronization with the spacecraft, instrument engineering and
safety, and science support.  The software is synchronized with the 1
pulse-per-second (1PPS) signal from the spacecraft.  The spacecraft
sends a mission-elapsed time (MET) to each of the instruments for a
common time tag.  The one-second period is subdivided into twenty
50-ms time periods, which allow for scheduling of software processes
and managing the S/C timing requirements for command and telemetry.  

The software processes for science provide support in 1) managing the
setting of HVPS optical power supply levels during science-data
acquisition (during HVSCI mode, the ESA, RPA and DFL are set by means
of lookup tables), 2) starting and collecting of counter data from the
PCEM, SCEM, and coincidence electronics, and 3) performing
calculations onboard using the counter data and telemeter science-data
products.

In nominal operations, commands are required for the high-voltage ramp
after a power on, but once the ramp has been completed, a small number
of commands are used for initiating science acquisition that set a
configuration for acquisition, the deflection angle and the telemetry
rates.

The FSW outputs three types of telemetry: instrument state
information, which can be read and acted upon by the spacecraft
software for onboard anomaly recovery; engineering telemetry
(housekeeping, messages and memory dump); and science telemetry
(real-time, summary and histograms).  The telemetry is summarized in
Table VI.

3.6.2 Science-Data Collection

Science data are collected in the HVENG and HVSCI modes.  HVENG was
used extensively during commissioning for initial HV ramp-up. HVSCI is
the primary SWAP science mode.  In HVSCI, the optical power supplies
are stepped every 0.5 seconds and the PCEM, SCEM and coincidence
counter data are collected with each step.  During each 0.5-second
period, approximately 100 milliseconds are allowed for the optical
power supply settling time and 390 milliseconds are allocated to the
counter sample period.  The duration of the latter is set with digital
hardware.  An overall cadence of 64 seconds consisting of 128
0.5-second steps defines the 64-second science-acquisition frames and
hence all science FSW activities. Two different methods of sweeping
during the 64 seconds can be performed - called the coarse-fine and
coarse-coarse sweeps - and are user selectable.  A typical coarse-fine
sweep consists of a 32-second coarse sweep, which covers the entire
energy range with 64 logarithmically-spaced optical power supply
values, followed by a 32-second (also 64 0.5-second steps) fine sweep.
A coarse-coarse sweep consists of two 32-second coarse sweeps
performed in one 64-second period.  For both sweep types, the optical
power supply values are read from one of several user-selectable
tables. 

For the coarse-fine sweep, the peak value of the coincidence counter
during the coarse sweep is calculated to determine the center of the
fine sweep so that finer resolution sweeping around the peak response
can be performed.

A 64-second acquisition frame period running a coarse-fine sweep from
SWAP calibration is shown in Figure 24, which plots the PCEM, SCEM and
coincidence count rates, and the ESA and RPA voltages.  A relationship
between energy and coincidence rate can be derived from such a plot
given that with each 0.5-second step, a different incident energy is
defined for each ESA-RPA voltage combination.  The figure shows how
the energy is swept from high to low, with the coarse sweep taking
place from time = 0 to 32 seconds and a fine sweep from time = 32 to
64 seconds.

Figure 24. A coarse and fine sweep acquired during SWAP ground
calibration activities.

To control acquisition while minimizing commanding to SWAP, science
data collection instructions are stored in user-selectable tables in
two identical EEPROMS.  The hierarchy of tables consists of:
- Plan (64 tables) - schedules which sweep table is active for each
  64-second acquisition frame.
- Sweep (16 tables) - defines which energy step and ESA/RPA
         table is used for each 0.5-second sample.  This determines,
         e.g., whether a sweep is coarse-fine or coarse-coarse.
- ESA (4 tables) - provides a mapping between energy step and
the digital-to-analog (DAC) value to program the ESA optical power
supply.
- DFL values (1 table) - maps energy step and angle to DFL DAC setting.
- RPA values (4 tables) - maps energy and angle to RPA DAC setting.
With the above tables, the user chooses a plan number and angle
setting to define science acquisition.

The summary is composed of a series of ongoing calculations based on
the fine portion of a coarse-fine sweep.  An array of coincidence
counts is collected and processed during each 64-second acquisition
frame.  Figure 25 shows an idealized coincidence count rate response
collected during a 32-second fine sweep in which the values related to
pseudo-velocity, pseudo-density and pseudo-temperature of a fine sweep
are defined.  For each of the summary values (angle, pseudo-density,
pseudo-velocity and pseudo-temperature), a running log of the sum,
minimum, maximum and variance are calculated.  While the sum, minimum,
and maximum are represented in telemetry as 32-bit numbers, the
variance is calculated by summing the squares of each of the values
and accumulated into a 64-bit number. Because the summary packet
defaults to transmitting once every hour, SWAP can use the summary to
collect data even during low-telemetry rate periods. The rate can be
altered via SWAP command for testing or other purposes.

Figure 25. Illustration of SWAP Summary Calculation

Figure 26 summarizes the histogram calculation.  The top half of the
figure shows a time line of two consecutive 64-second frames. The red
and blue plots show the coincidence counters during the coarse sweep
and the black plots show the response from the fine sweeps.  The red
plot shows that the peak was found at an energy index of 512, which
results in the FSW running a subsequent fine sweep from energy indices
480 to 543 in order to center the coarse peak in the middle of the
fine sweep.

Figure 26. Illustration of Histogram Calculation.  The bold vertical
lines on the upper plot show the maximum coincidence value for that
64-sample (32-second) span.

The coarse sweep data are shifted and accumulated into a normalized
histogram array that is composed of 2048 32-bit energy bins.  It is
"normalized" because each bin corresponds to a relative energy index
where 1024 always corresponds to the peak of the coarse array.  In
Figure 26, for Frame i, the fine peak energy index is 512.  This
number is used to calculate the shift required to move the
wider-spaced coarse sweep so its energy index 512 aligns with
histogram array index 1024.  Note that the wide spacing of the coarse
array is transferred also, so there are gaps between consecutive
coarse samples accumulated into the array.  In Frame i+1, the fine
peak energy index is 508, so a slight shift between the frames is
noted and results in the second coarse data set not aligning exactly
over the first coarse data set.

The number of times a particular bin (also with indices 0 through
2047) has been accumulated is telemetered along with the histogram
array so that the activity in each energy bin can be reported.  The
histogram packet defaults to being sent out approximately once per
day.  Because of the size of one histogram array, the packet is
divided and sent out over the course of 64 seconds.  The histogram is
used so that SWAP can collect data even during times where only
low-telemetry rates are allowed.  Like the summary packet, the rate
can be altered via SWAP command.

4.0 Instrument Performance

We characterized SWAP's performance through both laboratory testing
and electro-optics modeling. The purpose of the laboratory testing was
to calibrate the instrument and determine a ground truth.  The
computer models allowed us to examine more combinations of
incident-beam and instrument parameters than possible in laboratory
testing and to allow for analysis of other parameter combinations as
needed after launch.

4.1 Laboratory Testing

SWAP was extensively calibrated at the SwRI Ion Calibration
Facility. This facility produces ion beams from <500 eV to 51 keV. The
ion species can be selected for mass per charge from 1 to >40. Figure
27 shows the SWAP instrument mounted on the facility's 4-axis
positioning system, with the doors that protect the sensor for launch
in their open configuration.  Full functionality of the instrument was
demonstrated during calibration, including its ability to open the
doors while under high vacuum. The four-axis positioning system allows
vertical and horizontal displacement in the plane normal to the ion
bean and rotation about two orthogonal axes that cross in the beam
path: the outer rotation is about a vertical axis, and the inner,
nested, rotation about a horizontal axis.

Figure 27. SWAP instrument mounted on 4-axis positioning system 

The SWAP instrument was mounted with its center axis aligned to the
inner-stage rotation axis, which is equivalent to roll maneuvers the
New Horizons spacecraft will perform.  Because of SWAP's 17-cm
diameter aperture, outer-stage rotations resulted in translation in
the input aperture.  We used the horizontal stage to return the
aperture to its original position.  The positioning system's range of
motion is large enough to fully illuminate the instrument's FOV.

In addition to testing performed on individual components and on a
high-fidelity prototype, the full flight instrument underwent a number
of calibrations.  Calibrations were performed before and after
instrument level environmental testing, after refurbishment activities
following spacecraft environmental testing and a limited engineering
beam test immediately prior to delivery to the spacecraft for final
integration.  The instrument was calibrated over the full angular
response and energies from 500 eV to 8 keV.  Due to the nature of
electrostatic optics, the response of the instrument to other energies
can be easily scaled.  Because the solar wind protons have a nominal
energy of ~1 keV, a large fraction of the calibration data was taken
at this energy.  Results presented here are from the instrument in the
final delivered configuration.  

4.1.1 RPA

Figure 28 shows the measured energy response of the RPA to a beam of 1
keV protons.  The primary, secondary, and coincidence rates are shown
in blue, green and orange, respectively.  The RPA grids are curved,
rounding the cutoff of the transmission as compared to a planer RPA.
This is seen as the reduction of counts from 950 V to 1000 V. However,
as shown in Figure 10, by making small steps in the RPA voltage, we
are able to differentiate beams that are only ~1-2 eV apart at 1 keV.

Figure 28. Measured energy response of RPA to a beam of 1 keV protons

Figure 29 shows the general shape of the RPA response over the full
voltage range.  Normalized coincidence is plotted as a function of the
RPA voltage/beam energy. Because the individual RPA grids have a
finite thickness, they act as a series of electrostatic lenses.  The
features in this figure, including maximum near 0.945 of the 1 keV
beam energy, are due to the focusing properties, as discussed below in
Section 4.2.

Figure 29. General shape of RPA response over full voltage range

4.1.2 Deflector System (DFL) 

The FOV of the SWAP instrument is 276 degrees about the spacecraft roll axis
by 10 degrees in the plane normal to the roll axis.  SWAP uses an
electrostatic deflection (DFL) plate to increase the FOV out of plane
by up to 15 degrees in azimuth (alpha).  The required DFL voltage, normalized to
the beam energy goes as Equation 1, where alpha is the angle from the
plane.

For non-normal incident particles, the RPA response needs a
calibration correction. For an ideal RPA, non-normal incident
particles with E/q > (VRPA / cos(2 alpha)) would be passed.  The unique
focusing properties of the SWAP RPA lead to a more complicated
response function.  Figure 30 shows the normalized response of the RPA
as a function of the azimuth angle.  The RPA voltage is scaled by the
incident beam energy.  Protons with energy of 1000, 1010, and 1900 eV
are shown in red, green and blue, respectively.  We set the DFL
voltage according to Equation 1, leading to an empirically derived
calibration function for SWAP: Equation 2.

Figure 30. The normalized peak response of the RPA is shown as a
function of the incident azimuth angle. The RPA voltage is scaled by
the incident beam energy. Protons with energy of 1000, 1010, and 1900
eV are shown in red, green and blue, respectively. The DFL voltage was
set according to Equation 1. For an ideal RPA, non-normal incident
particles with E/q > (V_RPA*cos(2alpha)) would be passed. The unique focusing
properties of the SWAP RPA lead to a more complicated response
function, requiring an empirically derived calibration function as
shown in Equation 2.

For the SWAP RPA, non-normal incident particles with E/q > (VRPA /
f(alpha)) would be passed.

4.1.3 ESA

Figure 31 shows an example of the energy-angle response of the ESA to
a 1keV proton beam.  The RPA and deflector system's voltages were at
zero for the data shown here.  The ESA's analyzer constant is 1.88,
the energy resolution delta(E)/E is 8.5%, and the angular FOV, undeflected,
is 10 degrees.  

Figure 31. The energy-angle response of the SWAP instrument is shown
for a 1 keV proton beam.

In addition to performing the coarse energy measurement, the ESA also
blocks out UV light to the detectors.  A Krypton line source, with
emission at 123.6 nm and approximately the same intensity of the Sun
at 1 AU, was used to test SWAP for rejection of scattered light.  This
is a worst-case test since SWAP operates further out in the
heliosphere where the UV flux is lower by a factor of R2.  For all
angles tested with the UV source, the count rate never exceeded 1 Hz
for either primary or secondary detectors, and no coincidence events
were observed.

4.1.4 Detectors

Ions that make it through the sensor then pass through a thin carbon
foil and are measured by the PCEM.  Secondary electrons liberated from
the carbon foil are attracted to the SCEM.  Events that are measured
by both the PCEM and SCEM within 100 nanoseconds time are recorded as
coincidence (COIN) events.  

We determined the operation voltage for PCEM and SCEM by sweeping the
voltage while illuminating the instrument with a constant intensity 1
keV proton source.  Figure 32 shows the response of the PCEM, SCEM,
and the COIN rate as a function of voltage.  For the data shown here,
the voltage applied to the PCEM and SCEM were equal.  From these
tests, we selected a nominal operation voltage of 2100 V at the start
of mission, which is well into the saturation part of the CEM gain
curves.  The high voltage power supplies for the CEMs were designed to
output up to 4.5 kV, allowing us to increase the voltage to the CEMs
if their gain degrades over the mission lifetime.

Figure 32. PCEM (red), SCEM (gray), and coincidence (blue/green) rate
is shown as a function of CEM voltage. The magnitude of the PCEM
voltage is set equal to the SCEM voltage for the data shown here.

4.2 Instrument Modeling

We obtained instrument-modeling results with ion-optics simulations
using SIMION 7.0 software (Dahl, 2000). SIMION is an ion-optics
simulation program that models problems with 2D or 3D electrostatic
potential arrays. It creates potential arrays based on a user-given
geometry input (usually in the form of a text file containing
commands). The potentials of points outside the electrodes are
determined by solving the Laplace equation by finite difference
methods with a maximum of 50-million array elements. The time-step
integration for the calculation of ion trajectories uses a
fourth-order Runge-Kutta method, with ions flown individually or in
groups. Parameters such as position, velocity components, energy,
etc. can be recorded at different locations in the instrument model. 

The SWAP model consists of a 3D geometry with cylindrical
symmetry. The limitation of 50-million array elements prevents
reproducing the finer structure of the RPA design with the instrument
design in the same potential array. Therefore, the RPA grids are
modeled separately with a much finer resolution. We discuss the SWAP
model (instrument with "ideal" RPA) in Section 4.2.1, and describe the
RPA model in Section 4.2.2. 

4.2.1 Instrument Model with Ideal RPA

Our model geometry (Figure 33) fully conforms to the as-built
mechanical design of the flight instrument, with as much detail as the
resolution permits. This mechanical detail helps us account for
effects such as fringing fields. Figure 34 shows the focusing
properties of the ESA for normal incidence ions in two orthogonal
planes. In both planes, the ions are focused inside the ESA part and
then accelerated and re-focused as they reach closer to the
detectors. However, the impact positions on the primary CEM are nicely
spread in the funnel in order to avoid a hot spot. 

Figure 33. Cross section of the electro-optics model for the ESA. The
model has a main axis of symmetry. 

Figure 34. Focusing properties of the ESA shown in two orthogonal
projections. 

The instrument was thoroughly characterized with ion beam calibrations
(above) and the results agree very well with those from the
modeling. Table VII compares the energy-per-charge resolution, the k
factor and the angular acceptance. 

Table VII. Comparison of the instrument characterization between
measurements with ion beam and model

Figure 35 shows the angle-energy response. The angle and energy
binnings have been set to the same values for better comparison. The
right portion of the figure shows the normalized transmission, which
is color-coded using a logarithmic scale. Again, there is excellent
agreement between the modeling and the calibration measurements with
an ion beam.

Figure 35. Comparison between the angle-energy transmission function
of the instrument (left) and the model (right).

4.2.2 RPA Model

We modeled the RPA grids separately with much higher resolution,
matching the thickness and hole pattern of the mechanical
drawings. The model includes the curvature of the grids and drilled
holes, which are aligned for the four grids. This configuration does
not account for a possible misalignment of the hole patterns, which
may cause some differences between the model and actual
hardware. Figure 36 shows an isometric view of part of the RPA.

Figure 36. Isometric view of a part of the RPA. The voltage is applied
on the two middle grids. The holes are aligned so that their axes are
radial. The thickness of the grids and the diameter of the holes are
in accordance with the mechanical drawings. 

The ions are flown through the RPA grids, and the Cartesian
coordinates and velocity-vector components are recorded at the exit of
the RPA for those ions that make it through. Those ions are then flown
inside the instrument model described in Section 4.2.1, and an ion
that makes it to the carbon foil is considered to be detected.

In the following discussion, we present results using a 1 keV/q
parallel beam, which is uniformly distributed over the aperture of the
instrument at normal incidence (0 degrees in azimuth and elevation). Figure
37 shows a transmission curve as a function of the RPA voltage for an
ESA voltage of -524 V, with the ordinate in arbitrary units. The red
curve depicts the results of the RPA model and is compared to the
calibration data under the same conditions (black curve).

Figure 37. Relative transmission function as a function of the RPA
voltage for a 1 keV/e ion beam and an ESA voltage of -524 V. 

The RPA does not give a sharp cut-off at 1 kV because each hole of the
RPA grids acts like a tiny lens, focusing or defocusing the ion
trajectories. At a given E/q, the deflection increases with the RPA
voltage. Because of the grid thickness, some ions hit the inside of
the holes, which can also cause a non-uniform response below the
cut-off. Because the model holes are radially aligned, the E/q and the
RPA voltage have to be within a certain ratio for the trajectory to
deviate from a row of holes to another neighboring row. Therefore,
transmission varies as a function of RPA voltage for a particular
E/q. Although the general responses are similar, a discrepancy exists
between the calibration data and the simulations. The likely cause of
this discrepancy is that the RPA grid holes are not aligned along a
radial direction. In addition, the ion beam in calibration entered the
instrument with a small angle with respect to the normal of the RPA
grids (offset in azimuth and/or elevation angle) and the ion beam had
a small divergence (~1.0 degrees at 1 keV/e). These results highlight the
importance of using the flight calibration data in concert with the
instrument model for understanding the detailed response function.

We presented above the simulation results for one particular case,
however the ESA, deflector, and RPA responses were simulated with
numerous different initial conditions. The model aids in understanding
the instrument response and provides additional interesting
information that can be used for the analysis and interpretation of
the data. 

5.0 Science Operations and Data

This section describes the expected fluxes and predicted count rates,
instrument operations, and data pipeline. We used the fluxes measured
from previous missions and the instrument model described above to
estimate the anticipated count rates in the solar wind and at
Jupiter. The subsequent operations discussion describes the voltage
settings, which we vary according to the anticipated count
rates. Finally, the data pipeline section describes the data products
and level 2 data analysis. 

5.1 Fluxes and Predicted Rates

We predict SWAP count rates in the solar wind scaling from fluxes
observed by Voyager 2 over a large range of heliocentric
distances. The proton flux decreases with distance from the Sun
(Figure 38) primarily because the density decreases with distance. For
our calculations we use a constant speed of 1 keV (~440 km s-1). Fit
functions for the density, and density standard deviations are listed
below; the standard deviation for the speed is based on Figure 6 of
Richardson et al., (1996). Equation 3.

Figure 38. Flux of solar wind protons versus heliocentric distance.
The solid curve shows average values, while the other two show +/-1 sigma
probabilities.

We used the solar wind fluxes from Figure 38 and the instrument model
described above to estimate primary and secondary CEM count rates in
the solar wind as a function of heliocentric distance. The results are
shown in Figure 39. The instrument model includes the instrument
effective area and the ESA and RPA response curves. The model also
includes the deflector and angular response function and assumes a 1
keV beam centered in the ESA passband. The red line at 1 MHz indicates
a reasonable maximum count rate for the detectors. Although the count
rates are too high in the inner heliosphere with these voltage
settings, the RPA and ESA voltages can be adjusted to reduce passband.

Figure 39. Estimated count rates on the primary (top) and secondary
(bottom) CEM as a function of heliocentric distance for the fluxes
given in Figure 38.

As described above, the ESA and RPA can be crossed to provide a
reduced energy passband. Figure 40 shows the estimated count rates as
a function of the center energy of the ESA (kVESA) passband for a 1
keV beam with a flux that is one standard deviation higher than the
average flux at 5 AU. Each curve corresponds to sweeping the RPA and
ESA voltages while keeping the RPA voltage to ESA energy ratio
fixed. At a ratio of 1.03 the expected count rates at 5AU are reduced
to near maximum allowable levels, and for a ratio of 1.05 the expected
rates are <105 Hz.  For the last phase of commissioning we used a
ratio of 1.05 for safety.

Figure 40. Estimated count rates on the secondary CEM as a function of
the center of the ESA bandpass. Each curve corresponds to sweeping ESA
and RPA voltages such that the ratio of the RPA voltage to the center
energy of the ESA bandpass (VRPA/EESA) is fixed. The flux used is
higher than normal for 5 AU.

5.2 Instrument Operations

For the Jupiter phase of the mission, we determined that the radiation
levels during the encounter should be low enough to allow SWAP to
operate for the entire encounter. SWAP's shielding is similar to
Ulysses-SWOOPS, and the New Horizons spacecraft does not get as close
to Jupiter's radiation belts as Ulysses did. Phillips et al., (1993)
showed that Ulysses/SWOOPS operated as close as 16 R_J, compared to the
New Horizon's closest approach of about 32 R_J. Inbound to Jupiter, we
anticipate that real-time science mode data will be taken twice per
hour, with 20 measurements recorded per hour in Jupiter's tail. The
Jupiter tail observations continue until about 100 days after closest
approach, which corresponds to about 2200 RJ downstream. While not
designed for measuring Jovian ion distributions, we are hopeful that
some good new observations will be possible. 

Past Jupiter, we plan cruise phase science to study how the solar wind
and pickup ions evolve with radial distance from the Sun. For cruise
phase, real-time data will be recorded at a low rate and summary and
histogram data at a moderate rate. In addition, the gain on the CEMs
will be measured during each annual check out.

For the New Horizons primary mission, we will use SWAP's maximum
sensitivity and operate at its highest rate throughout the Pluto and
KBO flybys. These observations give a spatial resolution of <3000 km,
adequate to determine if Pluto has a bow shock and if so, locate it
with sufficient accuracy to make a good measure of the atmospheric
escape rate for the likely orientation of the IMF at flyby.

5.3 Data Pipeline

SWAP data are provided at three levels in the data pipeline. Level 1
is the raw decommuted data. Level 2 involves converting the level 1
data to more physical units, correlating data with thruster firings to
look for any enhanced backgrounds, calculating some orbital and
pointing information, and producing E/q spectrograms. For level 3, we
further process the data to obtain speed, temperature, and relative
density.

5.3.1 Level 1

Level 1 data are comprised of real-time science, summary, histogram,
and housekeeping packets. We decommute and place the raw data into
FITS files. Most of the SWAP data consist of binary tables with
columns for given instrument parameters and rows for each measurement
time. Since the histogram data consist of arrays, it is stored as an
image in the FITS file.  

We place all the real-time data for a given day into one file. A
complete real-time observation consists of 64 real-time packets. In
order to make sure that an instrument cycle is not split across two
daily files, a parameter in the real-time packet is used, which
indicates the beginning and end of the 64-second instrument cycle. We
also create daily files for summary data, but the histogram data are
handled differently since the histogram arrays are divided into 64
packets. A set of 64 histogram packets are combined and placed into
one file. 

We divide each FITS file into extensions. The first two real-time data
extensions contain header information and the real-time data as a
binary table. Additional table extensions contain housekeeping and
thruster-firing data. We use this same organization for the summary
data. The first two histogram extensions contain the array (image)
with the number of samples placed into each bin, and another image
holds the array of counts detected in each bin. The housekeeping,
thruster, source sequence count (used to check for data gaps) and
checksum (used to check for corrupted data) information is contained
in additional extensions. For real-time and summary files, the source
sequence and checksum information is not stored as separate
extensions, but directly inside the tables.

5.3.2 Level 2

We also store SWAP level 2 data as binary tables and images in FITS
files. The header information from level 1 is duplicated into level 2,
since the header contains level 1 processing information and
orbit-attitude information for the start time of the each file. The
second extension for the real-time and summary level 2 files contains
the corresponding data from those packets converted from raw to
engineering units, e.g., converting the RPA DAC step number to
voltage. The real-time files contain several additional extensions
that hold spectrogram information. These extensions consist of the
coarse scan (1st 32 steps) 2D spectrogram count arrays, the
corresponding 2D error spectrograms, the y-axis labels (energy label)
for the spectrograms, and the x-axis time label information. The
real-time and summary files have housekeeping, quality, and thruster
firing extensions. The histogram level 2 file format mirrors histogram
level 1 format, but includes quality flags. 

We base the flags that assess the quality of the data on housekeeping
data, orbit and attitude, and how well the software is
working. Currently, these flags are based on standard operating ranges
for housekeeping data, but additional ones can be created later as
needed.

We are developing SPICE code to calculate the time in UTC, position
and velocity of the spacecraft in a variety of coordinate systems,
field of view of the instrument, and the Sun location in instrument
coordinates. The boresight of SWAP is nearly aligned with the +Ysc
axis; therefore the SWAP frame kernel consists of only two small
rotations from spacecraft to instrument coordinates. We are
integrating SPICE code into level 2, and an orbit and attitude table
extension with all SPICE calculations for each measurement time is
planned for the level 2 files. Flags will be created to indicate
whether or not the Sun is in SWAP's FOV, since accurate solar wind
properties can only be determined if the bulk of the solar wind beam
enters the instrument. 

We will produce energy-time spectrograms from real-time science mode
data. These count rate spectrograms provide a way to examine the data
at high time resolution over the full energy range of the
instrument. Such spectrograms have proven useful for analyzing a wide
variety of distribution types. They have been used to analyze the
solar wind, comets, plasma boundaries, and shocks. As an example,
Figure 41 shows a spectrogram used to analyze ion measurements for the
comet Borrelly flyby. In this figure it is easy to locate the bow
shock and see the slowing of the solar wind due to mass loading.

Figure 41. Comet Borrelly flyby E/q-time spectrogram (Figure 1 of
Young et al. (2004) 

5.3.3 Level 3

Speed, temperature, and relative density depend heavily on the
calibration, orbit, and attitude information. We have planned these
data products for level 3 since the energy distribution (spectrograms)
and the orbit and attitude information, both of which are accomplished
in level 2 analysis, must be determined first.  Speed, temperature,
and relative density from the energy distribution/spectrum for the
solar wind measurements will be derived during solar wind
intervals. When only pickup ions are observed, the energy
distribution/spectrum is a more appropriate data product. We will use
ESA and RPA response functions, angular response function, instrument
solid angle, detector efficiencies, and effective area determined
during calibration to establish plasma properties and distributions
from the detected count rates. The effective area (Table II) is used
for a cold beam such as the solar wind when the entire solar wind beam
is in the field of view.  Note that the effective area is for a normal
incident beam and the transmission varies as a function of angle;
therefore, all of the angular dependent calibration functions
discussed in section 4 are also required to determine the
flux. Precise speeds and temperatures will be obtained from the fine
RPA scans. Lab calibration results already indicate the ability to
accurately determine the solar wind beam energy and speed with the RPA
fine scans. For times when the deflector is operated, we will use the
deflector's calibration data to analyze the data.

6.0 Early SWAP Observations

New Horizons was launched toward Jupiter and on to Pluto on January
19, 2006. Early operations included commissioning of all of the
instruments on board. The SWAP cover doors were opened, low and high
voltages turned on, and all aspects of the SWAP instrument were
demonstrated to be functioning nominally in space. Because of the
closeness to the Sun and high fluxes possible (Figures 38 and 39), we
ran SWAP with a very high crossing ratio of 1.05% between the ESA and
RPA. This crossing effectively limits the counts to ~1% of what would
be otherwise measured for all E/q values less that 2 keV/q; above that
value, the RPA is fixed at 2 kV, so the energy passband of the ESA is
effectively fully open.

Figure 42 shows an initial color-coded spectrogram of the
background-subtracted coincidence measurements of solar wind ions as a
function of E/q for January 8-20, 2006, when New Horizons was ~4.9 AU
from the Sun (~0.4 AU upstream from Jupiter). The lower trace is
produced by solar wind protons, while the upper trace comes from the
alpha particles (He++), which travel at approximately the same speed
as the protons as thus have twice the E/q. Alpha particle densities in
the solar wind are typically only ~4% that of the protons, so SWAP's
enhanced sensitivity above 2 keV/q is nearly ideal for making high
sensitivity measurements of these ions while still limiting counts due
to the protons to safe levels. Solar wind speed is a function E/q,
with 1 keV protons corresponding to typical, ~440 km/s solar wind with
larger (smaller) E/q representing faster (slower) wind speeds.

Figure 42. Early observations from SWAP in the solar wind at a
heliocentric distance of ~4.9 AU. The upper (lower) trace shows solar
wind alpha particles (protons) with the sensitivity significantly
reduced below 2 keV/q for safety reasons (see text). Times when the
proton and alpha particle beams are greatly reduced or disappear
correspond to times when the instrument was temporarily off or the
spacecraft pointed SWAP's aperture away from the Sun. The observed
stream structure and forward reverse shock pairs is consistent with a
CIR in the solar wind at these heliocentric distances.

A forward interplanetary shock passed over the New Horizons spacecraft
at ~1800 on January 11 and a reverse shock passed by at ~1300 on
January 14; the speed immediately following the latter shock was in
excess of 600 km/s. Such forward-reverse shock pairs, followed by
rarefaction regions (slowly decreasing speeds and consequent falling
E/q of the proton and alpha beams) are typical of corotation
interaction regions (CIRs) in the solar wind at these distances
[Hundhausen and Gosling, 1976; Smith and Wolfe, 1976]. These and the
continuing upstream measurements of the solar wind are critical for
understanding the external plasma environment as New Horizons entered
the Jovian magnetosphere late in February 2007.

7.0 Conclusions

The SWAP instrument on New Horizons will provide the best measurement
of the global atmospheric escape rate from Pluto. This escape rate
determines the type of interaction that Pluto has with the solar
wind. For a large escape rate, with many ions leaving the atmosphere,
the interaction should be comet-like, where ionization of atmospheric
material mass loads the solar wind, causing slowing of the overall
bulk flow of the solar wind and draping of the IMF. For lower escape
rates, the interaction is likely more like the interaction that the
solar and wind has with Venus.  There, the interaction size is much
smaller and the solar wind interacts directly with ions closely bound
to the planetary atmosphere, causing a much smaller region of draped
magnetic field. In either case, SWAP should also be able to identify
and locate a bow shock, if one exists.

In addition to measuring the overall escape rate from Pluto's
atmosphere and allowing us to develop an understanding of the solar
wind interaction with Pluto for the first time, SWAP should also be
able to measure hydrogen pickup ions around Pluto. In addition, in the
unlikely event that Pluto has any significant internal or remnant
magnetic field, SWAP should measure the effects that such a field
would have on the solar wind plasma interaction. Finally, en route to
Pluto and Kuiper belt objects beyond, SWAP should make unique
observation of the variation of heliospheric pickup hydrogen ions and
sample the Jovian magnetosphere.

The SWAP instrument was designed to meet the unique and challenging
requirements of measuring the solar wind and its interaction with
Pluto out at 30 AU, from a spacecraft that scans its orientation over
large angles in order to image Pluto and Charon with body-mounted
cameras as it flies by. In addition, because observations of the
plasma interaction is not primary science for the mission, resources
for SWAP were highly constrained and we had to remain focused on
developing a simple instrument optimized only for measurements of the
solar wind interaction with Pluto, while working within the detailed
constraints of the New Horizons spacecraft and mission plan.

To satisfy all of these requirements and constraints, we designed the
SWAP instrument to be cylindrical symmetric over 270 degrees of rotation
angle, and combined an RPA that acts a high pass filter, a deflector
ring, and an ESA.  Different combinations of the RPA and ESA produce a
variable passband and deconvolution of measurements made while
micro-stepping the RPA allows the measurement of very small changes in
the solar wind speed. Ions that pass through the SWAP electro-optics
are then measured in a coincidence detector section that both provides
redundancy for making the measurements and suppresses the background
generated by anything other than real ions transiting through the
instrument. The SWAP instrument has been well characterized by both
ground calibrations and a detailed instrument model. SWAP has been
successfully turned on in space and is making routine science
observations.

SWAP will make unique and exciting measurements en route to, at Pluto,
and beyond. SWAP will measure the slowing of the solar wind flow, mass
loading of the solar wind plasma, the escape rate of atmospheric
material from Pluto and the fascinating solar wind interaction with
Pluto for the first time.

Acknowledgments

We gratefully acknowledge all of the scientists, engineers,
technicians, and support personnel who have made the SWAP instrument
and New Horizons mission a reality! While it is not possible to list
everyone who has contributed to the SWAP instrument, a partial list
includes:  Maureen Ahr, Barry Eggington, George Geleta, Pat Gonzales,
Evan Guy, Helen Hart, Marc Johnson, Matt Maple, Greg Palacios, Robert
Rendon, Syrrel Rogillio, Carlos Urdiales, Kenny Whitton, and Mike
Young. We also gratefully acknowledge the work of Wendy Mills, who
helped assemble and edit this paper and two reviewers and Alan Stern,
who provided helpful comments on this manuscript. This work was
carried out under the New Horizons mission, which is part of NASA's
New Frontiers program. 

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