**Solar Wind Electron Proton Alpha Monitor**

**(SWEPAM)**

**for the Advanced Composition Explorer**

D.J. McComas, S.J. Bame, P. Barker, W.C. Feldman, J.L. Phillips, P.
Riley

*Space and Atmospheric Sciences Group*

*Los Alamos National Laboratory*

*Los Alamos, NM 87545*

J.W. Griffee

*Sandia National Laboratory*

*Albuquerque, NM 87185*

Submitted to

*Space Science Reviews*

*Special Issue on the ACE Mission*

14 April 1997

**Abstract**

The Solar Wind Electron Proton Alpha Monitor (SWEPAM) experiment
provides the bulk solar wind observations for the Advanced Composition
Explorer (ACE). These observations provide the context for elemental and
isotopic composition measurements made on ACE as well as allowing the
direct examination of numerous solar wind phenomena such as coronal mass
ejections, interplanetary shocks, and solar wind fine structure, with
advanced, 3-D plasma instrumentation. They also provide an ideal data
set for both heliospheric and magnetospheric multi-spacecraft studies
where they can be used in conjunction with other, simultaneous
observations from spacecraft such as Ulysses. The SWEPAM observations
are made simultaneously with independent electron and ion instruments.
In order to save costs for the ACE project, we recycled the flight
spares from the joint NASA/ESA Ulysses mission. Both instruments have
undergone selective refurbishment as well as modernization and
modifications required to meet the ACE mission and spacecraft
accommodation requirements. Both incorporate electrostatic analyzers
whose fan-shaped fields of view sweep out all pertinent look directions
as the spacecraft spins. Enhancements in the SWEPAM instruments from
their original form as Ulysses spare instruments include 1) a factor of
16 increase in the counting interval (and hence sensitivity) for high
energy, halo electrons; 2) halving of the effective ion-detecting CEM
spacing from \~5° on Ulysses to \~2.5° for ACE; and 3) the inclusion of
a 20° conical swath of enhanced sensitivity coverage in order to measure
suprathermal ions outside of the solar wind beam. New control
electronics and programming provide for 64 s resolution of the full
electron and ion distribution functions and cull out a subset of these
observations for continuous real-time telemetry for space weather
purposes.

**1. Introduction**

The Advanced Composition Explorer (ACE) mission was developed to examine
both the isotopic and elemental composition of our galaxy from the four
reservoirs of matter that are accessible to a direct measurement from
near the Earth. The first two of these reservoirs connect back to the
Sun through the solar wind and through solar energetic particles, the
third traces to the local interstellar medium by way of locally ionized
interstellar neutrals which form singly charged ions that are then
picked up by the solar wind, and the fourth traces to our galaxy at
large by way of galactic cosmic rays. These observations are carried out
with a suite of six experiments, which jointly measure ion composition
over the energy range from \~100 eV to several hundred MeV per nucleon,
with sensitivities that are one to two decades larger than previously
achieved. In addition, ACE carries three experiments that monitor the
interplanetary medium, thereby providing the context for the advanced
ion-composition experiments aboard ACE. These experiments are the Solar
Wind Electron Proton Alpha Monitor (SWEPAM), the energetic Electron,
Proton and Alpha particle Monitor (EPAM), and two magnetometers (MAG).

SWEPAM plays a critical role in ACE by providing the solar wind
observations that enable the understanding of the ACE composition
measurements in the context of previous knowledge about the solar wind.
This knowledge has been developed over more than three decades of solar
wind research. Specifically, SWEPAM will place the ACE composition
measurements in the context of low-speed streamer-belt flows, the
high-speed solar wind from coronal holes, coronal mass ejections (CMEs),
the various types and strengths of interplanetary shocks, magnetic
connection to the Earth's bow shock, and other solar wind structures.
Table 1, from the ACE project's Level 1 requirements definition, lists
the scientific objectives of the ACE mission. SWEPAM is considered a
primary measurement (P) for six of these objectives and an important
contributing measurement (C) for another 5. Thus, SWEPAM measurements of
the solar wind provide one of the critical cornerstones for the ACE
mission.

The SWEPAM experiment provides state-of-the-art measurements of electron
and ion distribution functions in three dimensions over all of the
velocity space needed to characterize the bulk flow and kinetic
properties of the solar wind. SWEPAM utilizes the refurbished flight
spare instruments from the Solar Wind Over the Poles of the Sun (SWOOPS)
experiment (Bame et al., 1992) onboard the joint NASA/ESA Ulysses
mission. Ulysses was launched on October 6, 1990. It was deflected out
of the ecliptic plane into its present high inclination polar orbit
about the Sun, using a gravitational assist at Jupiter, in late February
1992. SWOOPS has operated flawlessly since launch, thereby allowing
exploration of the solar wind as a function of heliolatitude for the
first time. Figure 1 summarizes the SWOOPS solar wind speed observations
as a function of latitude in a polar plot format. Note that while the
high-latitude solar wind is a comparatively steady flow of high speed
wind from both polar coronal holes, the low-latitude wind where ACE will
reside is a complicated mixture of slow flows from the streamer belt and
higher speed coronal hole flows.

The SWOOPS experiment comprises two separate instrument units, one for
electrons and one for ions. These instruments, although different in
size and detail, are based on the same basic principles. Both are
electrostatic energy per charge (E/q) analyzers followed by sets of
channel electron multiplier (CEM) sensors, which allow individual
particles transmitted through the analyzers to be counted. In both
cases, full complements of low and high voltage power supplies, memory,
data processing capabilities, and spacecraft interfaces are provided as
a part of the package such that each of the two instrument units is
operated completely independently of the other. For the SWEPAM
experiment on the ACE mission we have selectively modified, refurbished,
and improved the SWOOPS flight spare instruments. In particular, SWEPAM
provides enhanced sensitivity to suprathermal electrons (E \> 100 eV) to
allow a greater sensitivity for detecting CMEs and for establishing
their magnetic topology and evolution in interplanetary space. SWEPAM
also provides higher angular resolution measurements of solar wind ions
as well as increased sensitivity to suprathermal ions coming from
directions adjacent to the bulk solar wind in order to provide a bridge
between the energetic ions detected using the suite of ACE composition
experiments, and the solar wind thermal population.

This paper provides the primary documentation source for the ACE/SWEPAM
experiment . In Section 2 we briefly review solar wind science and place
the anticipated SWEPAM observations and scientific results in the
context of previous observations. Sections 3 and 4 provide detailed
descriptions of the SWEPAM ion and electron instruments, respectively.
In each we: 1) summarize the electro-optical properties and general
specifications of the sensor, 2) show the mechanical design of the
sensor, 3) discuss the electronics associated with each instrument, and
4) show the results of sensor calibrations. While sufficient detail is
given for the interested reader to understand the implemented mechanical
and electronics designs, these two sections are arranged so that the
detailed subsections can be skipped without losing an understanding of
the top-level design. Section 5 describes the commanding, data handling,
modes of operation, and the real-time solar wind capability of SWEPAM.

**2. Scientific Objectives**

The primary purpose served by SWEPAM is to provide the solar wind
context for the other experiments on ACE. Connection to global
heliospheric structure will be facilitated by correlation of SWEPAM data
with simultaneous measurements made using the SWOOPS experiment on
Ulysses. During the prime portion of the ACE Mission, Ulysses will be
positioned near 5.4 AU, close to the ecliptic plane. The heliographic
locations of ACE relative to Ulysses between 1997 and 2002 are shown in
Figure 2. Times of near radial lineups are indicated by the
intersections between the dotted and solid traces in the top panel.
These lineups occur in early 1998, 1999 and 2000.

It will also be critical to place the new ACE observations in the
context of synoptic changes in solar wind structure that are associated
with phases of the solar cycle. The time period of the ACE mission will
extend from near solar minimum toward the coming solar maximum. Over
this portion of a solar cycle, the photospheric and coronal magnetic
fields evolve from simple, dipolar-like configurations, with axes nearly
perpendicular to the ecliptic plane, to configurations of considerable
complexity. This complexity is characterized by a rapid spatial
evolution of the photospheric magnetic field and accompanying time
variability, first associated with the appearance of magnetic active
regions at relatively high heliographic latitudes, and later to their
appearance at progressively more equatorial latitudes. Understanding ACE
observations as a function of the solar activity cycle will be
facilitated by comparison of these observations with measurements made
from the particle and fields experiments onboard ISEE-3 in the late
1970s and early 1980s. These observations were made from the same
Sun-Earth L-1 neighborhood to be occupied by ACE, but in a previous
solar cycle. Specific contributions by SWEPAM in support of a number of
the major scientific thrusts of the ACE mission are presented below.

**2.1 Formation of the Solar Corona and Acceleration of the Solar Wind**

Synoptic studies of solar and solar wind structure during the 1970s
revealed an evolving global configuration that followed the pattern of
changes in photospheric and coronal magnetism mentioned above. Near
solar minimum, the global solar wind consists of high-speed flows
emanating from both solar poles separated by a roughly planar, low-speed
streamer belt at equatorial latitudes that is nearly parallel to the
ecliptic plane (see Fig. 20d of Hundhausen, 1977). A similar global
structure was observed by Ulysses during its fast-latitude scan from
-80.2o to +80.2o heliographic latitude between September, 1994 and
August, 1995, as shown on the left-hand side of Figure 1. This interval
occurred just prior to the minimum of solar activity during the present
activity cycle. An expanded version of this traversal, including the
observed proton density scaled to 1 AU, is shown in Figure 3 (assembled
by B.E. Goldstein using SWOOPS plasma data and a Yohkoh X ray image of
the corona, kindly provided by the Yohkoh mission of ISAS, Japan).
Ulysses observed nearly structure-free, high-speed solar wind (V \> 700
km s-1) for all heliographic latitudes poleward of about 22o, and
variable, but generally low-speed solar wind equatorward of 22o. This
high-speed - variable-speed transition was recently reconfirmed by
Ulysses as it approached aphelion within about 29o of the ecliptic plane
(upper right quadrant of Figure 1).

If these conditions prevail throughout the early portion of the ACE
mission, as expected, then phenomena associated with the streamer belt
and its embedded heliospheric current sheet will dominate observations
made by all experiments onboard ACE (see the yellow stripe in the bottom
panel of Figure 3). Observations made by ISEE-3 during the same phase of
the solar cycle between 1978 and 1980, and by Ulysses just after its
launch early in the current solar cycle, suggest that ACE will also
likely spend a fraction of its time in those portions of coronal
hole-associated high-speed streams that are adjacent to streamer belt
flows. Transitions between streamer belt and coronal hole associated
flows correspond to the flow-speed variability encountered by Ulysses at
intermediate heliographic latitudes shown in Figures 1 and 3. The
correspondence between the solar wind flow state observed by Ulysses and
the structure of the corona as observed in X rays by Yohkoh, is shown by
comparing the panels in Figure 3. The high speed wind is clearly
associated with the large polar coronal holes. While most, if not all,
of the solar wind must expand along open field lines, the low-speed wind
seems to be associated with the bright, largely closed magnetic
structures. The intermediate and variable solar wind appears to come
from with the margins between these two regions.

During the approach to solar maximum, an ever-increasing number of
magnetically active regions rise through the photosphere. These then
interact both with the preexisting dipolar magnetic field and with each
other, producing a series of rapid reconfigurations that contribute to a
pattern of solar magnetism characterized by a regular (and at times,
extreme) time variability. These reconfigurations lead to an
increasingly complex magnetic configuration. The global structure of the
corona and resultant solar wind change accordingly in response to these
underlying changes. In particular, an ever greater fraction of the solar
wind encountered far from the Sun is associated with coronal mass
ejections (CMEs) and other transient flows. In addition, some of the
apparently time-stationary high-speed flows encountered outside of the
streamer belt originate from relatively small unipolar photospheric
regions that evolve on timescales of months, in contrast to the
timescale of years over which the much larger area coronal holes evolve
near solar minimum.

Although much is already known about the evolution of the large-scale
pattern of solar wind structure in response to the solar activity cycle,
very little is known about the mechanisms that connect the flow state of
the slow wind with plasma conditions in the corona. Even more
fundamental, we do not yet know the mechanisms that heat the solar
corona and accelerate its plasma outward to form the solar wind. One of
the major scientific goals of SWEPAM is to provide the link between the
data from the composition experiments on ACE and the sources of the
various types of solar wind back in the solar corona.

**2.1.1 Streamer Belt**

Bright, high-density streamers are clearly seen during solar eclipses
and are prominent in coronagraph images of the Sun. These streamers,
which are also associated with the equatorial bright regions seen in the
soft X ray images of the Sun (e.g., Fig. 3), give rise to highly
structured but slow solar wind flows. These flows are characterized by
large variations in speed, density, temperature, heat flux, helium
abundance, and heavy ion charge-states (e.g., Feldman et al., 1977,
1981; Borrini et al., 1981). Some of these variations probably reflect
fine spatial-scale structure in the streamer belt owing, at least in
part, to the meandering of the neutral line in the photosphere. The
interplanetary extension of this neutral line is a current sheet that is
embedded in streamer belt-associated flows. In addition, variability in
the low-speed wind may be caused by quasi-continuous, fine-scale
magnetic restructuring at low altitudes above the photosphere (Axford,
1985). Some of this restructuring may be caused by the different rate of
rotation of fine-scale photospheric magnetic fields relative to that of
the deep seated currents that are responsible for large-scale patterns
of solar magnetism (Sheeley et al., 1975), and some may be due to the
emergence and decay of active regions within the streamer belt.

Attempts have been made to associate peculiarities in the internal state
of solar wind ions and electrons with each of the foregoing
possibilities, without definitive success. The problem in interpretation
stems from essential ambiguities in relating the plasma at 1 AU with
production mechanisms in the corona. A partial list of solar wind
observables includes: 1) double ion streams, possibly driven by magnetic
reconnection (Hammond et al., 1995; Feldman et al., 1996); 2) heat flux
dropouts, possibly signifying magnetic disconnection from the corona
(McComas et al., 1989); 3) heat flux enhancements, possibly caused by
transient-driven shocks that are magnetically connected to the
spacecraft, (Feldman and Marsch, 1997); 4) non-compressive density
enhancements (NCDEs) (Gosling et al., 1977); 5) the intensity and shape
of suprathermal electron velocity distributions, possibly signifying
major magnetic restructuring of the inner heliosphere; (Gosling et al.,
1987 ; Feldman and Marsch, 1997); and 6) the He to H relative abundance
variations, which may be caused by changes in the structure of the
chromosphere-corona transition zone (e.g., von Steiger and Geiss, 1989).
A major scientific objective of the SWEPAM experiment is to relate such
observables to peculiarities in isotopic and elemental abundances
measured by the composition experiments on ACE.

**2.1.2 Coronal Mass Ejections**

The interplanetary manifestations of CMEs in the solar wind are varied.
Perhaps the single best identifier is the occurrence of
counter-streaming suprathermal electrons (e.g., Gosling, 1990, and
references therein). The appearance of CMEs in coronagraph images shows
radially evolving sets of structured loops of enhanced density; these
structures are interpreted as tracing sets of nested magnetic loops that
are closed at both ends to the Sun. Connection to the hot corona at both
ends of these loops should, in turn, drive an outward-directed electron
heat flux along both coronal footpoints. The consequent field-aligned
streams of suprathermal electrons that carry this heat flux are thought
to be the counter-streaming electron signatures observed in
interplanetary space. An example of such a signature is shown in C and D
of Figure 5 (from McComas et al., 1994). By comparison, regions A and E
show more typical halo electron distributions where direct connection to
the hot corona is made in only one direction along the magnetic field.

Although the presence of counter-streaming suprathermal electrons is
highly likely, it is neither a necessary nor a sufficient interplanetary
signature of a CME. Many other plasma and magnetic field identifiers
have also been associated with CMEs in the solar wind. These identifiers
include: 1) a generally enhanced helium abundance that is often
characterized by large temporal and/or spatial variations (Hirshberg et
al., 1972); 2) unusually low ion and electron temperatures (Gosling et
al., 1973; Montgomery et al., 1974); 3) macroscopic rotations of the IMF
interpreted as a magnetic cloud (Klein and Burlaga, 1982); 4)
charge-states of heavy ions that indicate a coronal temperature larger
than about 2 x 106 K (Bame et al., 1979); and 5) counter-streaming
suprathermal ions (Marsden et al., 1987 ).

The solar wind identified with CMEs is often highly structured. Not only
do the various observables used to identify a CME often vary
considerably during any given event, they often do not coexist in any
given parcel of CME plasma. Indeed, sometimes they do not even overlap
spatially. These observations have led to a picture of solar wind CMEs
in the solar wind as consisting of a raisin-pudding conglomerate of
small-scale structures convecting together in the ambient interplanetary
medium. Figure 6 shows a schematic sketch of a CME in the solar wind
that combines features presented by Bame et al. (1979) and Gosling and
McComas (1987). This basic picture has been reinforced by recent
observations of a CME observed simultaneously at high and low
heliographic latitudes by Ulysses and IMP-8 (Gosling et al., 1995). It
is curious that while the He abundance in this CME was \>10% at times at
low latitude, in the same CME at high latitude, it remained constant at
the high-speed wind value of \~4.4% (Barraclough et al., 1996).

**2.1.3 Outstanding Questions**

A common thread that connects streamer-belt flows with CMEs is that they
both originate from regions that overlie complex magnetic topologies in
the photosphere that, at times, undergo large reconfigurations. The
streamer belt generally overlies most zones of high magnetic activity
that consist of complex, strong magnetic field configurations having a
high multipolarity and a relatively short lifetime (e.g., Hundhasuen,
1977). CMEs are thought to result from a large-scale restructuring of
longer-lived magnetic field structures (Low, 1996). Time sequence
observations of the corona in white light (see, e.g., Hundhausen, 1988),
EUV (Dere, 1994), and in X rays (Tsuneta, 1996) have graphically
illustrated the consequences of these reconfigurations on the structure
of the transition region and overlying corona. These range from the
ponderous uplift of large volumes of coronal gas (a CME) to the
injection of fine-scale, high-speed jets of relatively cold (H-α surges
and EUV explosive events, 104 to 105 K) and very hot (X ray jets, 106 to
107 K) gas into the preexisting, ambient solar wind.

All the foregoing phenomena and their consequent solar wind extensions
to interplanetary space are intrinsically time dependent. Although the
relatively gross aspects of their origin (free energy contained in the
photospheric magnetic field), and their effects on the low corona
(convected enthalpy and kinetic energy of plasma bulk motion), are known
from remote-sensing observations, not very much is known about the
mechanisms that drive them. It seems reasonable to presume that magnetic
reconnection in some form is the prime mechanism that effects the
transfer of energy from the pre-existing magnetic field to the plasma
and energetic particle population. However, we do not know the altitude
at which reconnection is initiated. Nor do we know the plasma properties
in regions that undergo reconnection, or the relative fraction of the
magnetic energy that goes into plasma bulk motion, ion heating, electron
heating, and the energetic particle population. Are magnetically-closed
structures ever produced by this reconnection, or do they always produce
complex flux rope topologies containing a mixture of open, closed, and
disconnected field lines (i.e., Gosling et al., 1995)? Schematic
examples of several possible magnetic topologies are shown in Figure 7
(from Phillips et al., 1992). What is the time scale of the initial
reconnection event? Are slow-mode shocks generated by these events as
suggested by Petschek (1964), and if so, what is contribution to the
plasma structure of the streamer belt? At what altitude are fast-mode
shocks initiated? Do the flux ropes continue to evolve as they expand
outward into the solar wind? Is there a component of the streamer-belt
flow at 1 AU that is not associated with solar activity (i.e., that is
driven by a quasi-static pressure gradient as first suggested by Parker
(1958))? If so, what is the heating mechanism that supports a
quasi-static, million-degree corona and where does this heating occur?

Answers to many of these questions have been offered based on analyses
of previous particle and fields observations. None of these, however, is
definitive because of essential ambiguities in interpretations of most
of the data that was available. Specifically, many characteristics of
these data (such as plasma ion and electron temperature and particle
energy spectra) are affected by so many processes that occur between the
Sun and Earth that they contain highly ambiguous information about
conditions near the coronal base where the solar wind is launched.
However, the ambiguity in interpretation of the isotopic, elemental, and
charge-state composition of plasma and energetic particles from the Sun
is considerably reduced because these parameters do not change greatly
above the coronal base during their transport to 1 AU. We therefore
expect that it should be possible to make considerable progress in
finding solid answers to many of the foregoing questions through a joint
analysis of the solar wind data returned by SWEPAM and the composition
measurements made by the comprehensive composition payload on ACE.

**2.2 Particle Acceleration and Transport**

Another major thrust of the ACE mission is to increase our understanding
of the origin and evolution of the energetic particle populations that
are found almost ubiquitously throughout the inner heliosphere. An
important component of this population is accelerated in interplanetary
space, and most of it must be transported across large distances in the
solar wind before being detected at 1 AU. Numerous analyses have shown
that the forward and reverse shocks that bound co-rotating interaction
regions (CIRs) beyond 1 AU are copious accelerators of particles. An
example of one such interaction region observed by Ulysses at 5 AU is
shown in Figure 8 (Gosling et al., 1993). Note the large enhancements in
both ion and electron suprathermal populations extending to considerable
distances upstream from both the forward shock (day 20 \~0500) and
reverse shock (day 22 \~0200) that bound the CIR. Other important sites
of particle acceleration occur near forward shocks that are driven
through the ambient plasma by fast CMEs.

The seed population for particle acceleration can originate in the
thermal and suprathermal components of the solar wind, or in the
population of pick-up ions that are injected into the solar wind from
the neutral component of the interstellar medium, which interpenetrates
the inner heliosphere. This latter injection results from 1)
charge-exchange collisions with solar wind ions; 2) photo-ionization by
solar EUV photons; and 3) impact ionization by solar wind electrons.
Although charge-exchange collisions and photo-ionization generally
dominate impact ionization in the injection of interstellar neutrals
into the solar wind, shock-heated electrons upstream of some CMEs can
sometimes make important (and for helium, dominant) contributions, as
illustrated in Figure 9 (Isenberg and Feldman, 1995; Feldman et al.,
1996).

Although the shapes of particle energy distributions largely result from
processes that occur in the neighborhood of shocks, they evolve in
important and characteristic ways while streaming from these shocks to
their point of detection. This evolution is controlled both by the
topology of the IMF that is carried by the solar wind and by the nature
of the waves that fill interplanetary space. The effect of waves on
energetic particles is enhanced if the magnetic topology is that of a
complex flux rope or if the imbedded magnetic field intersects a shock
front at more than one location. In such cases, the time available for
scattering a seed particle in transit through a given interplanetary
distance is longer and there is more time for interactions.

**2.2.1 Corotating Interaction Regions**

Shocks driven by corotating interaction regions do not generally form at
radial distances less than about 2 to 3 AU, and so will only
occasionally be directly detected by ACE near 1 AU. SWEPAM data can,
however, be used to predict CIR onsets and their locations relative to
Earth using appropriate MHD models. Predictions of these models can be
checked using the solar wind observations made by the SWOOPS instrument
onboard Ulysses, which will be located just beyond 5 AU near the
ecliptic plane throughout the prime portion of the ACE mission (Figure
2). This combination of dual-point observations and MHD model
predictions should provide the needed context for composition
observations of backstreaming populations of energetic ions accelerated
by the shock pairs that bound CIRs at large heliocentric distances.

**2.2.2 Forward Shocks**

Forward shocks, driven by fast CMEs, accelerate significant numbers of
thermal and suprathermal ions and electrons to high energies (see, e.g.,
Figure 8). These shocks, and the associated enhanced populations of
suprathermal ions and electrons, will be detected and fully
characterized by SWEPAM. Special attention was given in upgrading the
physical configuration of the SWEPAM ion analyzer, and in implementing
data accumulation cycles for both plasma ion and electron experiments,
to return fully three-dimensional energy-angle spectra having sufficient
velocity-space resolution and counting statistics to provide a bridge
between the solar wind thermal population and the energy spectra of
energetic particles. The shapes of these distributions should be
sufficiently detailed to enable definition of the magnetic topology of
the ambient plasma through which the shocks propagate. This topology
will be of fundamental importance in assessing the type and efficiency
of acceleration processes at work at these shocks.

**2.2.3 Pick-up Ions of Interstellar Origin**

SWEPAM data will also be used to support studies of the solar wind
pickup of newly ionized interstellar neutral atoms. Most important for
the pickup of interstellar hydrogen will be SWEPAM's measurement of the
solar wind mass flux because the ionization rate is dominated by
charge-exchange collisions with solar wind protons.

Past studies of ISEE-3 and Ulysses electron data show that just upstream
and downstream of strong interplanetary shocks, electron impact
ionization can enhance ionization and hence pickup of interstellar
hydrogen atoms by roughly 10% (Isenberg and Feldman, 1995; Feldman et
al., 1996). These same measurements also indicate that electron impact
ionization can dominate photo-ionization for interstellar helium atoms
(see, e.g., Figure 9). These effects are doubly significant because the
subsequent acceleration of these newly injected ions to higher energy is
enhanced in the neighborhood of shocks. SWEPAM data can also be used to
help in studying particle transport through the solar wind by providing
knowledge of the magnetic topology in the neighborhood of these shocks.
This knowledge will result from study of the shapes and intensity of
suprathermal electron and ion velocity distributions.

**3. Solar Wind Ion Instrument (SWEPAM-I)**

The solar wind ion instrument, SWEPAM-I, contains a spherical section
electrostatic analyzer (ESA). Ions enter the sensor through its aperture
and are immediately introduced into the gap between the electrostatic
analyzer plates. A negative high voltage on the inner plate of the
electrostatic analyzer biases the analyzer such that only ions within a
narrow energy per charge range (\~5%) and azimuthal angular range
(3°-4.5°, depending on polar angle) are permitted to pass through the
analyzer. The transmitted ions are detected in the channel electron
multipliers (CEMs) located immediately behind the analyzers. The
aperture of the sensor is aligned such that its fan-shaped field-of-view
rotates about the spacecraft spin axis. Figure 10 schematically displays
the SWEPAM-I pointing and electro-optics. We define two angles through
the normal to the sensor aperture: the polar angle, θ, in the plane of
this fan (plane of the right hand sketch), where 0° is parallel to the
sun-pointing spin axis direction of the ACE spacecraft. The azimuthal
angle, φ, is in the plane perpendicular to the fan (plane of the left
hand sketch).

The normal to SWEPAM-I's aperture is pointed 18.75° in polar angle away
from the sunward pointing spin axis direction on the ACE spacecraft.
This allows SWEPAM-I to measure ions arriving at polar angles from 0°
(along the sun-pointing spin axis) out to \~65°. Ions in the main solar
wind beam, which usually arrive from polar angles of \<25°, are measured
by CEMs on both sides of the fan. The offset angle of the sensor's
pointing will interleave observations from one half of the rotation with
those from the other half with a consequent improvement in polar angle
resolution from 5° down to 2.5°, if the spacecraft spin axis is
perpendicular to the Sun-facing deck of the ACE spacecraft. While the
ACE design goal is for the spin axis to be perpendicular to this deck to
within 0.5°, the exact orientation will not be known until after
deployment of the solar panels on orbit.

Ions entering the ESA at different polar angles with respect to the
aperture normal end up being counted in different CEMs, thus identifying
the incident polar angle of the ion in the plane of the acceptance fan
of the instrument. As ACE spins, the CEMs sweep out small conical
segments of the sky, centered on the spacecraft spin axis direction. The
combination of CEM identification and spin phase of the spacecraft
provide unique directional information about all ions measured in the
instrument. Further, the high voltage step of the ESA selects the E/q of
transmitted ions. This combination of CEM number, spin phase, and ESA
step level, which is equivalent to polar angle, azimuthal angle, and
energy of the incident ions, respectively, provides the basic matrix of
measurements in which SWEPAM-I counts are collected.

SWEPAM-I instrument capabilities are summarized in Table 2. This
instrument makes full 3-D plasma measurements of protons and alpha
particles every 64 s. In its search mode (typically once every 32
minutes) SWEPAM-I covers the entire 260 eV/q - 36 keV/q energy range.
Between search modes, SWEPAM-I tracks the solar wind and measures every
other of 40 energy levels around the peak; track mode provides 5% energy
resolution measurements every 64 s. In addition, observations from two
adjacent track modes can be interleaved to give 2.5% energy resolution
with 128 s cadence. Section 5 describes the modes of operation in more
detail.

Figure 11 displays a photograph of the SWEPAM-I instrument which mounts
on the sunward-facing instrument platform of the ACE spacecraft. The
instrument is comprised of two major subassemblies, a large cylindrical
"sensor" assembly, and a rectangular "electronics box". The two are
joined by angled mounting brackets, which set the pointing angle of the
sensor head and, hence, instrument aperture with respect to the
spacecraft spin axis.

**3.1 SWEPAM-I Sensor**

Components of the ion sensor are housed in the ribbed, thin-walled,
drum-shaped container visible in Figure 11. All of the parts are
attached to the end covers of the drum. The curved-plate electrostatic
analyzer is mounted on the inside of the sunward drum face, and views
space through a single, oversized entrance aperture. The ESA is composed
of nested spherical section plates having a 105° bending angle.
Conically-shaped stiffening ribs were machined as integral parts of the
curved aluminum alloy plates in order to maintain their dimensional
stability. Both of the plates were coated with copper and blackened
using an Ebanol-C process in order to reduce UV scattering into the
sensor. The electrostatic analyzer has a gap width of 2.84 mm between
plates and an average radius of 100 mm.

Sixteen channel electron multiplier detectors, or CEMs, are arrayed such
that their 7 mm funnels have 5° polar angle separations around the edge
of the gap where the particles exit the analyzer plates. Each is
individually mounted on a ceramic card which also carries glass-coated
thin film resistors and bare ceramic coupling capacitors needed to bias
the anode slightly positive with respect to the exit end of the CEM, and
to couple out the signal. This arrangement maintains cleanliness in
order to guarantee long CEM lifetimes (e.g., McComas and Bame, 1984 and
references therein), permits replacement of individual CEMs as needed,
and enables the precise spacing that is required for the fine angular
resolution of this instrument.

In the twin instrument on the continuing Ulysses mission, a
motor-actuated aperture wheel provides the variable sensitivity required
to make measurements across the wide range of heliocentric distances
covered by that mission. For ACE, which remains at a fixed distance of
0.99 AU, this mechanism was replaced with a single, oversized entrance
aperture wide enough to expose the entire width of the analyzer gap, and
extending 7.8 mm along the gap. The aperture is covered with a simple,
spring-loaded dust cover. Redundant pyro-activated dimple motors actuate
the release mechanism by command, on orbit.

The ACE mission limits the Sun-spacecraft spin axis angle to between 4°
and 20°, far less than the Ulysses mission. Therefore, we were able to
build new instrument mounting brackets, visible in Figure 11, that are
configured so that, when the instrument is mounted on the spacecraft,
rotation of CEMs 1-12 about the spin axis covers the total angular range
of solar wind ions expected throughout the mission. The central axis of
the measurement array, parallel to the spacecraft spin axis, is located
one quarter of the way between the centroids of CEMs 5 and 6 (θ =
18.75°).

The new mounting scheme also makes it possible to devote four of the
SWEPAM-I polar angle detectors to high-sensitivity measurements of
suprathermal ions. We therefore added a thin aluminum exit aperture mask
covered with 70 lines per inch, 92% transmissive nickel mesh interposed
between the exit of the analyzer plates and the CEMs. This mesh, biased
at the same voltage as the funnels, precludes electric fields from the
ESA from reaching into the funnels and affecting the CEM counting
efficiencies. This mask also provides two different sizes of exit
apertures. Analyzed ions enter funnels 1-12 through apertures tailored
for solar wind ion fluxes near 1 AU. These apertures are 0.4 mm wide,
selecting ions exiting only from the very center of the gap. CEM funnels
13-16, which will be exposed only to suprathermal ions and not to the
solar wind ion beam because of their angular positions, have circular
apertures in the exit mask large enough to transmit ions from across the
ESA gap into the CEM funnels.

Each of the sixteen CEMs is connected to its own dedicated
amplifier-discriminator circuit on its own miniature printed circuit
board, located on the outside of the sensor box. Signal lines are fully
shielded, only a few centimeters long, and pass through the SWEPAM box
wall on individual miniature feedthroughs, in order to avoid any
cross-talk or noise pickup at the inputs to these sensitive circuits.
These circuits amplify and discriminate pulses from the CEMs; each
discriminator has two, selectable level settings of 1x10^6^ and 2x10^7^
electrons per pulse. The lower, more sensitive setting is used for all
science observations while a combination of alternating between these
levels and stepping the CEM high voltage is used for periodic CEM gain
calibrations. Normalized pulses from the 16 circuits are cabled to
counters in the electronics box.

Since the pumping speed through the aperture is small, SWEPAM-I is
fitted with an additional pump-out baffle box to enhance the evacuation
of the interior of the instrument. The blackened baffle box requires a
minimum of seven surface reflections between its two open ends. This
design, fully proven on Ulysses, does not allow for any appreciable
background from photons or charged particles entering through the
pump-out channel. To preserve the cleanliness of the sensor interior,
the aperture and baffle openings were hermetically sealed during the
long storage period after Ulysses development was completed. For the
bulk of the ACE observatory testing the instrument has been continuously
purged.

**3.2 SWEPAM-I Electronics**

With the exception of the sensor assembly amplifier-discriminators
circuits described above, the ion instrument's electronics are housed in
the attached electronics box (Figure 11). The purpose of the support
electronics is to 1) furnish programmed high voltages for the analyzer
which sets the ion energy per charge passband, 2) bias the CEMs at
commandable high voltages appropriate to achieve high and stable
counting efficiency, 3) condition and provide low voltage power for the
various electronic circuits and high voltage power supplies, and 4)
process and accumulate pulses to form a digitized data string for input
into the spacecraft command and data handling system. Figure 12 provides
a block diagram of the SWEPAM-I electronics.

In order to configure the Ulysses SWOOPS instruments for use on ACE, two
fundamental changes were required. First, the instrument's Ulysses
interfaces needed to be adapted to the ACE spacecraft. Second, the
flight processing capabilities needed to be increased in order to meet
the ACE science objectives. In order to achieve these objectives, we
took advantage of the instrument's modular design and replaced a select
number of logic modules and the back plane (motherboard), leaving the
heritage high voltage and low voltage supplies untouched.

The electronics box has four cavities: the low voltage power supply,
high voltage power supply, motherboard and main electronics cavities.
The motherboard cavity is the primary wiring cavity, interconnecting the
spacecraft interface connectors, power supplies, processing electronics
and sensor electronics. All signals enter and exit this cavity on
connectors. The main electronics cavity (tall, central portion of the
electronics box) contains five boards or modules: the processor
(PRO-11), counters (SPM-14), signal buffer/level-shifter module
(BUF-19), spacecraft interface (SIM-11), and high voltage controller
(BAM-16) module.

The combined ion instrument has four external connectors: spacecraft
command and data handling, spacecraft power, pyro for cover actuation,
and a test connector. The first three are mated into the spacecraft wire
harness while the latter is used for test purposes only. SWEPAM-I has
two grounds. Signal ground is isolated from the electronics box housing
but connected to the electrically isolated sensor head. The signal
ground is carried back to the spacecraft single point ground on a wire
in the spacecraft harness. SWEPAM-I dissipates 3.1 watts when it is
fully functional at expected count rates. There are no survival heaters
but the instrument is thermally bonded to the thermally controlled
spacecraft structure. The instrument temperature is monitored by a
sensor mounted on the bulkhead between the motherboard and the low
voltage power supply.

**3.2.1 Power Supplies and Controller**

The SWEPAM-I electronics box connects directly to the spacecraft 28 V
power and contains both low and high voltage power supplies to produce
and condition voltages used in the instrument. Two boards make up the
low voltage power supply. The first is a DC/DC converter which isolates
the spacecraft primary power ground from the instrument's secondary
power ground. This heritage supply "chopper" was designed to be
synchronized to Ulysses' power system. The ACE spacecraft does not have
power system synchronization so this supply free runs at 55 +/-0.5 kHz.
The second board of the low voltage supply rectifies, regulates and
filters the outputs from the multi-tap isolation transformer. Spacecraft
interface signal line drivers and receivers are powered by +5 V; +/-15 V
is provided for the high voltage control module and used to condition
analog telemetry signals output to the spacecraft, +8 V is provided for
the sensor amplifier-discriminator circuits; the balance of the
processing electronics are powered by an 8 V to 5 V linear power
regulator on the spacecraft interface module.

Two encapsulated high voltage supplies are provided: a -2.4 kV to -3.9
kV CEM bias supply and a -15 V to -2 kV electrostatic analyzer field
supply. Both supplies are driven by the BAM-16 module. The high voltage
supply lines are shielded, coaxial cables; the supply return current is
carried on the shield. The CEM bias supply assembly includes the
feedback controlled drive circuit, drive current monitor, step-up
transformer, 5 stage diode/capacitor voltage multiplier and the high
voltage monitor/feedback resistor divider network which supplies the
reference level for the feedback controller. The ESA field supply
assembly includes the feedback controlled drive circuit, drive current
monitor, step-up transformer, 5 stage diode/capacitor voltage
multiplier, feedback controlled high voltage shunt regulator and a
regulator performance monitor. In order to improve efficiency, this
supply is designed to vary both the feedback controlled multiplier and
regulator references such that the multiplier output is 10% above the
regulator output. Most of this 10% head room is required to allow for
the multipliers' slow recovery rate in the "flyback" to start the next
sweep. It turned out that the CEM and ESA high voltage supplies did not
function well at the 55 kHz free running power supply chopper frequency.
Consequently, it was decided to allow these clocks to free run at their
natural frequency (57 kHz) so none of the power supplies are
synchronized to the spacecraft power system or to each other.

The BAM-16 high voltage control module provides the supplies with
switched input power, drive clocks, reference levels, monitor
conditioning and over current protection. Each supply is activated by
its own ON/OFF control signal; an over current condition will turn this
switch off requiring the ON/OFF control signals to be cycled to
reactivate the supplies. The processor selects the desired CEM bias
level with a four bit control word. The BAM-16 translates this into 16
equally spaced levels from -2.4 kV to -3.9 kV (100 V spacing). The
processor also selects the desired ESA field level with an eight bit
control word. The BAM-16 translates this into 256 logarithmically spaced
levels of which only the first 200 (-15 V to -2 kV) are used.

**3.2.2 Counters, Processor, and Interface Modules**

The SWOOPS digital electronics were all powered with 8 V. It was not
uncommon in the late 1970's and early 1980's when the SWOOPS instruments
were developed to find CMOS processors, memories and peripherals that
would operate at these or even higher levels. By the 1990's when the
retrofit of the SWOOPS to the SWEPAM was made, essentially all CMOS had
been standardized to run on 5 V. In order to reuse the heritage counter
module (SPM-14), a new module (BUF-19) was added to level shift the 8 V
amplifier-discriminator signals to 5 V level before routing to the
counter module. The SPM-14 module contains 16 counters, one per CEM, and
the bus interface circuits. The TA585 counters used are custom 16 bit
counters which interface directly to an eight bit processor bus. While
these devices do not have a spill flag, the SWEPAM sample intervals and
electronic count rates preclude an overflow.

The Ulysses 1802 microprocessor-based processor module and two memory
modules were replaced with a single new processor module, PRO-11. PRO-11
includes an 80C51 microcontroller, 7.4 MHz crystal controlled clock, 2K
Bootstrap PROM, 28K EEPROM , 48K RAM, watchdog timer, inflight
calibration pulser and an 8 channel power supply monitor multiplexer.
The timer, pulser, multiplexer controller, memory controller and I/O
decoding are all implemented in a single Field Programmable Gate Array
(FPGA).

Using lookup tables, its telemetry clock (not the spin clock) and its
own chip counter/timer circuit, the microcontroller generates a timing
function which sequentially disables the counter inputs, commands a new
ESA level, moves the data from the counters into RAM, clears the
counters, allows high voltage to settle then re-enables the counter
inputs to collect data at the new E/q level. Each instrument data mode
has a lookup table which defines the mode in terms of the E/q level and
count interval, high voltage settling times, number of E/q levels per
spin sector, number of spin sectors per spectral segment (roughly
equivalent to a spin) and the number of spectral segments per mode.

Bootstrap PROMs contain the firmware which initialize the instrument
when the processor is first activated and provide the facilities to load
all flight software into EEPROM or RAM. These devices are read-only and
contain two copies of the bootstrap codes to ensure there is always a
way to correct or tune the flight software. The EEPROMs contain the bulk
of the flight software. The instrument will be launched with the most
current version of the flight software resident in the EEPROM. If any
software uploads are required after launch, they will most likely be
limited to revisions of the data mode lookup tables that are moved to
RAM by the startup codes.

To accommodate these instruments on the ACE spacecraft, a new spacecraft
interface module was designed to replace the Ulysses heritage interface.
This module, the SIM-11, includes the command receiver, framing pulse
receiver, data transmitter, high voltage control latches (CEM and ESA
ON/OFF, CEM Level and ESA Level) and an 8 V to 5 V linear power
regulator to supply the processor electronics. In order to keep the
number of processor interrupts to a minimum, the data transmitter uses a
first-in/first-out (FIFO) memory device permitting the processor to
buffer an entire minor frame's worth of data between readouts. Except
for this FIFO and the linear regulator, all functions on this module are
implemented in a single FPGA.

**3.3 SWEPAM-I Calibration**

Calibration of SWEPAM-I was performed at the Los Alamos plasma analyzer
calibration facility. The angular and energy response functions for all
sixteen CEMs were measured using a known, unidirectional, 5 kV proton
beam. In order to maintain absolute cleanliness for the CEMs during the
calibration, we chose to mount only the cylindrical sensor head, and not
the electronics box, within the vacuum chamber. The sensor head was
mounted on a computer-controlled three-axis table (two orthogonal
rotation stages and one translation stage) which allowed ions to be
directed at the aperture over all possible input angles for the
instrument. A 3-D array of sensor transmission/ counting efficiencies
was obtained as a function of polar angle, azimuth angle, and ESA
voltage (giving E/q) for each CEM. The response function of the
instrument is completely defined by these sixteen 3-D arrays.

An example of a 2-D cut through the 3-D response function is shown for
CEM 4 in Figure 13. Levels of transmitted ions from an ion beam arriving
at a fixed 20° polar angle are color-coded as a function of azimuthal
angle and ESA voltage. The azimuthal and energy responses are coupled as
expected for a spherical section electrostatic analyzer (Gosling et al.,
1978); the polar response is essentially decoupled from these two.
Figure 14 shows three 1-D cuts through the response function for channel
4 at the centroids of each of the other two parameters. In each, the
measured number of counts (solid line) is fit with a theoretical curve
(dotted line). In Figure 14 (a) the distribution of counts for CEM 4 is
shown as a function of analyzer plate voltage for the 5 keV proton beam
in a look direction given by φ=0° and θ=25.8°. The response is
approximated by a gaussian profile. The azimuthal profile of the CEM at
an analyzer plate voltage of 293 V and polar angle of θ=25.8° is also
approximated by a gaussian profile and is shown in Figure 14 (b); .
While the responses as a function of both analyzer plate voltage and
azimuth are approximated by a gaussian, the response in polar angle is
best approximated by a trapezoid. Comparison of the measured counts with
the best-fit trapezoid is shown in Figure 14 (c) for an azimuthal angle
of 0° and analyzer plate voltage of 293 V. The measured analyzer
constant for SWEPAM-I, given by the ratio of incident particle energy to
plate voltage , is \~17.1.

Calibration data similar to those shown for CEM 4 were taken for each of
the 16 CEMs in SWEPAM-I. Figure 15 shows a polar cut through all 16
channels at 0° azimuth. This cut appears to be more variable than one
might have expected, however, multiple runs confirmed the variations
shown. In addition, the roughness of these curves is at least in part
attributable to the fact that the central azimuth was slightly different
(\<1°) for each channel and that the very narrow angular response
functions were cut in slightly different places. In any case, the full
response function information is used in characterizing each channel's
calibration curves.

Once these 16 arrays were fully measured, the polar, azimuthal, and
energy responses were characterized and the geometric factors for each
CEM were calculated. Geometric factors were calculated by numerical
integration of the three response functions (energy, θ-angle, and
φ-angle) over all energies, and all angles. These geometric factors
provide the necessary conversion from the measured counts to more
physically meaningful distribution functions, from which one can
calculate the moments of the distribution such as density, flow velocity
vector, temperature matrix, and the heat flux. Table 4 summarizes the
calibrated geometric factors for SWEPAM-I.

**4. Solar Wind Electron Instrument (SWEPAM-E)**

Like SWEPAM-I, the electro-optical design of the solar wind electron
instrument, SWEPAM-E, is based around a spherical section electrostatic
analyzer. For SWEPAM-E, the bending angle of the analyzer is 120°.
Figure 16 schematically displays the SWEPAM-E electro-optics. Electrons
enter the instrument through its aperture which is pointed normal to the
spacecraft spin axis. A positive high voltage on the inner plate of the
ESA biases the analyzer such that only electrons within a range of
energies (12%) and azimuthal angles (9˚ at θ = 0° to 28˚ at θ = +/- 70°)
are permitted to pass through the analyzer and are detected in the seven
channel electron multipliers (CEMs) located immediately behind the
analyzer.

Polar and azimuthal angles are defined just as they were for SWEPAM-I.
Electrons entering the ESA at different polar angles with respect to the
aperture normal end up being counted in one of the seven CEMs which
gives 21° polar angular resolution for electrons entering in the
wide-azimuthal angle, fan-shaped FOV. The plane of symmetry of the fan
is parallel to the spacecraft spin axis and the middle CEM (#4) viewing
angle is centered on a plane parallel to the spacecraft equator (θ =
90°). As ACE spins, the SWEPAM-E fan-shaped field of view sweeps out
\>95% of 4π sr, missing only small (\~10° half angle) conical holes
centered along directions parallel and anti-parallel to the spin axis.
Electrons with energies from 1.6 eV - 1.35 keV are measured in 20
contiguous energy bins determined by instrument's high voltage stepping.
Again, the combination of CEM number, spin phase, and ESA step level,
which is equivalent to polar angle, azimuthal angle, and energy of the
incident electrons, respectively, provides the basic SWEPAM-E
measurement matrix. The general SWEPAM-E instrument parameters are
summarized in Table 2.

Figure 17 displays a photograph of the SWEPAM-E instrument which also
mounts on the sunward-facing instrument platform of the ACE spacecraft.
Like SWEPAM-I, this instrument is comprised of two major subassemblies,
a ribbed, thin-walled, cylindrical sensor assembly, and a roughly
rectangular electronics box. The sensor cylinder mounts directly onto an
extension of the electronics box such that the instrument aperture
points at right angles to the spacecraft spin axis.

**4.1 SWEPAM-E Sensor**

The curved-plate electrostatic analyzer is mounted on the inside of the
outward pointing drum face, and views space through a single,
rectangular entrance aperture. The analyzer is composed of nested 120°
spherical section plates of aluminum alloy that are grooved in addition
to being blackened in order to minimize backgrounds caused by UV
scattering through the plates and photoelectron and secondary electron
production in the plates themselves. The electrostatic analyzer plate
system has a gap width of 3.5 mm and an average radius of 41.9 mm. Seven
large-funnel (11 mm diameter) CEMs are arrayed with 21° separations
along the analyzer plate's exit. A thin aluminum exit aperture mask
covered with 70 lines per inch, 92% transmissive nickel mesh (not shown
in the simplified drawing) is again interposed between the analyzer exit
and CEM funnels to preclude electric fields from affecting the CEM
counting efficiencies. Each CEM is mounted on its own ceramic card which
also carries the resistors and coupling capacitors needed to bias the
anodes slightly positive with respect to the exit of the CEMs and to
couple out the signals. Just as for SWEPAM-I, this arrangement maintains
cleanliness, permits replacement of individual CEMs as needed, and
enables the precise spacing that is required for the fine angular
resolution of this instrument. Each of the seven CEMs is connected to
its own dedicated amplifier-discriminator circuit; these circuit boards
are identical to those used for SWEPAM-I.

Electrons enter the analyzer through an entrance aperture which is wider
than the analyzer gap and runs 10 mm along it. At launch, this aperture
is also covered with a spring-loaded dust cover to protect the CEMs from
contamination. Redundant pyro-activated dimple motors release a
mechanism by command on orbit, and allow the cover to be retracted from
over the aperture. A blackened pump-out baffle box, similar to the one
used on SWEPAM-I, is provided to enhance the evacuation of the interior
of the instrument. Throughout most of the observatory level testing
SWEPAM-E has also been under continuous purge.

Because of the sensitivity of low-energy electrons to spacecraft surface
charges, the SWEPAM-E requires a special spacecraft thermal blanket.
This blanket maintains the thermal characteristics of the spacecraft, is
conductive to carry off any surface charge build-up to the spacecraft
bus, and also produces no magnetic field. The side of the SWEPAM-E cover
containing the aperture is made with an outer surface of indium/tin
oxide (ITO) coated kapton. This material has the same thermal properties
as the regular spacecraft blanket (which is made of beta cloth - a
fiberglass material with a teflon coating). A short (\~25 cm) vertical
section of the blanket covering the spacecraft panel below the SWEPAM-E
aperture is also made of this ITO coated material. Finally, the original
(non-conductive) thermal material on the spacecraft radiator near
SWEPAM-E was removed and replaced by a conductive ITO coated
silver/teflon film.

**4.2 SWEPAM-E Electronics**

The electron instrument electronics perform the same functions and have
a significant amount of commonality with those of the ion instrument. In
particular, we determined that it was most cost effective to use a
single design for the new electronics hardware for the two instruments.
Since the ion instrument's processing requirements were more demanding,
they drove the processing and memory resources built into SWEPAM-E as
well as SWEPAM-I. SWEPAM-E dissipates 2.7 watts when it is fully
functional at expected count rates. Its thermal interface is the same as
the ion instrument's interface: no heaters, one temperature sensor, and
hard mounted to the spacecraft. Figure 18 provides a block diagram of
the SWEPAM-E electronics

This electronics box houses the high and low voltage power supplies,
motherboard, wiring, and main electronics. The main electronics cavity
contains five boards or modules: the processor (PRO-11), counters
(SPM-17), signal buffer/level-shifter module (BUF-19), spacecraft
interface (SIM-11), and high voltage controller (BAM-25) module. Of
these, only the SPM-17 and BAM-25 are unique to the electron instrument.
The rest of the boards are identical to those used in the ion instrument
and described above. SWEPAM-E has the very same four connector
interfaces as the ion instrument right down to the pin out and signal
ground arrangement with one exception. Since the pyro connector is
mounted to the sensor housing rather than the electronics box, its shell
is grounded to signal rather than chassis ground. Special precautions
were taken to assure that this arrangement did not cause a loss of
isolation between the grounds.

While the SWEPAM-E heritage low voltage power supplies are of a
different design than those of SWEPAM-I (the SWOOPS ion sensor placed
additional requirements on the supply for its aperture stepper motor),
they are functionally equivalent, including the same 28 V isolation,
free running chopper frequency and secondary powers and grounds. The two
encapsulated high voltage supplies are a +2.4 kV to +3.9 kV CEM bias
supply and a +8.6 V to +300 V electrostatic analyzer field supply. Each
supply interfaces identically with the BAM-25 module and sensor head as
it did in SWEPAM-I. The CEM bias supply is identical to that of the ion
instrument except for output polarity. The ESA field supply assembly is
very much like that of the ion instrument except for its voltage range
and number of levels. The SWEPAM-E ESA supply output is 32 logarithmic
levels from +8.6 V to +300 V. The first eight levels are not used for
the ACE mission. Conversely, in order to extend the top-end range of
electron science observations, SWEPAM-E uses the top two levels which
were not used by the SWOOPS.

The BAM-25 is the complement of the BAM-16 in the ion instrument and
functionally differs only in the number of ESA step levels it controls
and its references levels. In order to get the full range out of this
ESA supply, the 1 V reference had to be set slightly low (\~0.9 V) so
that the supply would not sag at its highest level. This also decreased
the output voltage for all other levels proportionately. Consequently,
the SWEPAM-E ESA levels are offset from those on SWOOPS. The SPM-17
counter module is the equivalent to the SPM-14 module in the ion
instrument except that it has only seven counters instead of 16. Since
the electron instrument did not need as much RAM memory for data
buffering, the second RAM chip was omitted during assembly.

**4.3 SWEPAM-E Calibration**

Calibration of the electron analyzer was performed at the Los Alamos
plasma instrument calibration facility in a similar manner to that of
the ion instrument. Because our facility produces ions rather than
electrons, we chose to also calibrate SWEPAM-E using a 1.05 keV proton
beam. This was made possible simply by using a negative high voltage on
the inner analyzer plate and using ions with energies well above the
several hundred volt post-acceleration bias on the CEM funnels. This
procedure also has some advantages over using electrons as the ions are
far less affected by the Earth's magnetic field than are electrons.

The calibration procedure was essentially the same as for the ion
instrument. A 2-D (energy/azimuthal) angle cut through the response
function for CEM 5, at its polar angle centroid (-21°), is shown in
Figure 19. Figure 20 displays the central 1-D polar, azimuthal, and
energy cuts for this channel in the same format as Figure 14. The
distribution of counts for Channel 5 is shown as a function of analyzer
plate voltage for 1.05 keV protons in a look direction given by φ=0° and
θ=-21° in (a); the central azimuthal and polar response are shown in (b)
and (c). Again, the responses in (a) and (b) are approximated by simple
gaussians. For electrons, the response in (c) is approximated by a
skewed gaussian. Finally, Figure 21 shows the polar cut at 0° azimuth
through all seven CEMs in SWEPAM-E. The ratio of electron energy to ESA
voltage, or analyzer constant, for SWEPAM-E is \~4.3. The
channel-specific geometric factors were derived by the same method as
was done for SWEPAM-I and are given in Table 5.

Since the SWEPAM instruments, and particularly SWEPAM-E, bear a
considerable resemblance to the SWOOPS instruments onboard Ulysses,
comparison of calibration results obtained from SWOOPS prior to launch
have been made with the SWEPAM calibration; overall the comparison is
excellent. As an example, in Figure 22 we compare the geometric factors
calculated for the ACE plasma instruments with the geometric factors
calculated for Ulysses. With the exception of CEMs 3 and 4, the
efficiencies are very close. The excellent overall quantitative
agreement between these curves, measured in different decades in
different calibration facilities, provides support for the accuracy of
both sets of measurements. Since the ACE instruments are essentially
made up of Ulysses spare parts, the flattening of the geometric factor
at CEMs 3 and 4 is probably due to differences in the sensor hardware
and/or CEMs. Nevertheless, such differences are minor and should not
impact the science.

**5. Instrument Operations and Data Handling**

Owing to the large differences between the ULYSSES and ACE orbits, as
well as the large leaps forward in computing capability over the years
between the two instruments' development phases, the SWEPAM data modes
and operations have been completely reworked. As such, this section
describes a completely new development for the ACE program. Several of
the features of the new operations allow SWEPAM-I and SWEPAM-E to
provide substantially improved data products over their Ulysses
counterparts.

Throughout this section we will use the term "cycle" to describe the
top-level instrument process. Cycles are comprised of sequences of
"modes" which are further comprised of sequences of "segments" which
occur over intervals that are roughly equivalent to the spacecraft spin
period. The nominal ACE spin rate is 5 RPM or 12 s per spin. Operation
of SWEPAM is synchronized with the telemetry rate. This choice yields a
nominal data accumulation period (or segment) of 12.8 s which is also
just slightly longer than the longest possible spacecraft spin period
forecast for the mission (12.2 s). The difference between the actual
spin period (which will not be known until after launch, and may change
slightly over the course of the mission) and 12.8 s causes the data
sampling window to precess around from one spin to the next. As will be
described below, this is not a concern because the SWEPAM microprocessor
selects out the most interesting samples in the full data matrix for
telemetering.

Table 6 summarizes all science data sets for both SWEPAM-I and SWEPAM-E;
details are given in the following subsections. The first five products
are derived from the ion instrument and the last three are derived from
the electron instrument. The data mode from which the data set were
derived is shown in column 2. Column 3 provides a brief description of
the data set. The fourth column describes the data set matrix in terms
of energy (E), polar angle or CEM (θ), and azimuth or spin angle (φ).
The notation θ−φ is used to indicate that a mask has been applied and
only a subset of θ−φ angles are telemetered. The last column provides
the time resolution and cadence for each data product. Thus all science
data sets have a basic time resolution of 64 s. The cadences indicated
are based on expected data mode cycles, but can be easily changed during
the mission.

**5.1 Commanding and Housekeeping**

Once the SWEPAM instruments have been powered by the spacecraft, they
are operated by two distinct classes of ground commands; hardware
commands (hard commands) and software commands (soft commands). Both
classes are transmitted to the instruments from flight operations ground
computers through the spacecraft as serial data commands. The hard
commands are dispatched by a hardware decoder to control critical
hardware redundancies and interlocks. The soft commands, are dispatched
by one of two command processors in the flight software and can either
update the EEPROM flight codes and RAM data base or select and control
the data acquisition functions. All SWEPAM operational commands are
defined in the flight operations data base and are issued from command
procedure files which verify the commands and instrument responses. The
spacecraft has provisions for stored and autonomous instrument commands.
There are currently no plans to use stored commands to operate the
SWEPAM instruments. The only autonomously implemented commands for
SWEPAM are power shedding if either a spacecraft or instrument power
fault is detected.

The SWEPAM instruments each have three housekeeping "channels". The
analog telemetry consists of eight low voltage monitors multiplexed on
one of the three analog signals and 11 high voltage control monitors
multiplexed on two other analog signals. These signals are converted by
the spacecraft and provided at a rate of one sample per major frame
(i.e. 16 s). Eight major frames are required to collect a full set of
these monitors. In addition, housekeeping data are provided in the first
40 bits of each major frame to enable operation and monitoring of the
instrument by the flight operations team. Twenty-four bits of this same
header are output in the RTSW data link. The header is the key to data
set formats, identifying the data set type and frame count. Operating
parameters required to interpret the science data are embedded in the
science data sets, in data set headers, and spectral segment headers as
appropriate.

**5.2 Ion Modes and Operations**

SWEPAM-I has four commandable data modes and two commandable data
cycles. The data modes are; engineering (ENGI), calibration (CALI),
solar wind ion (SWI), and search supra thermal ion (SSTI). Data modes
can be run individually or as part of a data cycle. The data Cycles are;
normal and calibration. The normal cycle is a sequence of one SSTI
(search) mode followed, nominally, by 30 SWI (track) modes. The number
of SWI modes, and therefore the duration of this cycle, is command
adjustable. Once selected, the normal cycle runs continuously until a
command is sent to select a new data mode or cycle. The calibration
cycle consists of a sequence of two CALI modes to ensure that a change
in the solar wind does not corrupt the data set. This cycle will run
once after being commanded and then the operation returns to the data
cycle (or data mode if no cycle is running) that was in effect at the
time the calibration cycle command was received. Except in the event of
a "cycle abort" command, a cycle can not be interrupted. Therefore, a
newly commanded data mode or data cycle will not take effect until the
current cycle has completed.

With the exception of the ENGI data mode, which becomes a contingency
mode after instrument turn-on and initialization in space, all other
modes collect data for four major frames or 64 s. The data collected
during this period comprises one full spectrum subdivided into five 12.8
s spectral segments. Since the nominal spin period is 12 s and the worst
case predicted spin period is 12.2 s (a 2% error), each spectral segment
contains data from slightly more than one complete spin. For all data
modes, count accumulation begins at the start of a major frame, and
readout of a spectrum begins one major frame later.

Telemetry rate restrictions necessitate the use of several techniques to
reduce the number of bits actually transmitted. One of these techniques
is a compression algorithm that is used to compress the 16 bit counts
and count sums to eight bits. The algorithm, which makes use of a
compression table, introduces errors that are typically 1-2% and in the
worst case, no larger than 3%.

**5.2.1 Solar Wind Ions (SWI)**

The SWI data mode is often called the "Track Mode" since it uses the
level with the maximum counting rate (*L*mx) in the previous SWI or SSTI
data set to adjust the E/q range of the current collection. Azimuthal
angle samples are taken from each of the 16 CEMs (polar angles, θ) at 40
(8 per each of 5 spectral segments) consecutive even or odd E/q levels
starting with the level *L*mx -25. In a repeating sequence of SWI modes,
it is this subtraction of an odd number from the previous *L*mx that
results in the alternating even-numbered/odd-numbered starting energy
level. Since the E/q spacing between adjacent levels is 2.5%, every
other level corresponds to an E/q spacing of 5%. Thus it is possible to
interleave adjacent data sets and provide the highest possible
resolution of 2.5% (at the expense of a factor of two in temporal
resolution). There are 61 φ-angle sectors in a spectral segment and at
the nominal 5 rpm spin rate the φ-angle sectors at a given E/q level are
spaced by 6.19° and are 0.59° wide. Since there are 58 φ-angles per spin
at this nominal spin rate, the SWI mode intentionally over samples. The
sequence of voltage stepping is illustrated schematically in Figure 23.
Each horizontal profile corresponds to a spectral segment, which is
roughly equal to one spin.

The SWI mode produces three data sets. The first (DSWI) consists of
counts from CEMs 1 through 12. The second (DSTI) consists of counts from
the unmasked CEMs 13 through 16. The third (DRTSW) is a subset of the
DSWI data set used for real-time solar wind (see section 5.4).

For DSWI, the telemetry rate for ACE precludes the transmission of all
732 pixels (12 θ-angles × 61 φ-angle) for each of 40 E/q level, thus a
set of θ-φ angle masks have been defined that select those pixel
locations which well characterize the solar wind peak and adequately
cover the suprathermal wings of the distribution function. There are 12
masks, one for each θ-angle (CEM). The mask used for a given spectral
segment is determined by the CEM with the peak count rate during that
segment. Briefly, the masks were chosen by calculating the angular
offset of each pixel look direction and discarding those pixels with an
offset greater than 25° from the peak. Of the remaining pixels, 96 were
selected such that the pixels were more finely spaced for smaller
angular offsets. Thus the output consists of counts at each of 40 energy
levels for the 96 θ-φ angle positions which best define the solar wind
beam. Although this represents only 13% of the total number of pixels,
tests with simulated data sets (based on measurements by the ion plasma
instrument onboard Ulysses) indicate that even under extreme solar wind
conditions, where the ion beam is extremely broad (i.e. slow and hot),
almost none of the non-zero data pixels will end up being discarded by
this scheme. As noted above, the SWI mode collects more than one spin's
worth of data. The mask application allows for adequate ESA high voltage
settling time at the start of each spectral segment and then collects
only the first 58 of the 61 φ-angle sectors to preclude seeing the solar
wind beam twice in one data set. This limit is adjustable to allow for
variations in the spacecraft's spin rate.

Samples from the outer four θ-angles, representing the suprathermal
population from CEMs 13 through 16, are presented in the DSTI data set.
Again, telemetry restrictions preclude the transmission of the entire
array and two averaging procedures are performed onboard. First, the
eight E/q levels measured in each spectral segment are summed in pairs
to form four E/q levels per spectral segment. Second, the first 60 of
the 61 φ-angles per pair of E/q levels are further summed to form six
φ-angles per spectral segment. Thus the final DSTI data set contains 29
E/q levels, four θ-angles, and six φ-angles.

**5.2.2 Search/Suprathermal Ions (SSTI)**

During the SSTI mode, a fixed voltage stepping scheme is used with a
step interval of 5 at low energies (first two spectral segments) and 4
at higher energies (last three segments). This mode allows construction
of a complete suprathermal ion array and enables a robust determination
of the solar wind energy peak, *L*mx, in the event that the peak
tracking is lost during a SWI mode. The resulting energy range covered
by these steps ranges from \~500 eV/q to 35.2 keV/q. At low energies
(segments 1 and 2), the E/q spacing is 12.5% and at higher energies
(segments 3-5) the E/q spacing is 10%.

Two data sets are produced from the samples collected in the SSTI mode;
DSSTI and DSTI2. These are exactly the same in format as DSWI and DSTI,
respectively, although the energy range covered is broader and the
energy resolution is less. Thus the DSSTI contains a complete but coarse
energy spectrum from 500 eV/q to 35.2 keV/q, while DSSTI provides a
spectra at about one quarter of the energy resolution of the DSWI mode,
but over a much broader range of energies.

**5.2.3 Ion Calibration (CALI)**

The CALI mode allows the determination of the gain saturation level of
the CEMs. The E/q level is fixed at the peak of the distribution (*L*mx
level) from the previous search or track mode. The CALI mode data
collection consist of 16 CEMs × 5 CEM bias levels × 108 φ-angles × 2
threshold levels. The discriminator threshold is alternated between A
and B in successive counting intervals of the same φ-angle, with the
transition occurring as the counters are read. The operating level of
the CEM bias supply (CEMLev) is set by command. During the CALI mode the
actual CEM bias level is varied about this operating level for each of
the five spectral segments in this sequence; CEMLev−2, CEMLev−1, CEMLev,
CEMLev+1 and CEMLev+1.

The data set produced by the CALI mode consists of counts from all 16
CEMs. Again bandwidth restrictions require some averaging of the data.
The telemetered data set is created by forming separate sums for the A
levels and B levels from four consecutive φ-angle sectors. Thus, this
data set consists of 16 CEMs × 5 CEM HV levels × 2 threshold levels × 
27 φ-angle sums.

**5.3 Electron Modes and Operations**

SWEPAM-E has five commandable data modes and two commandable data
cycles. The data modes are; engineering (ENGE), calibration (CALE),
normal solar wind electron (NSWE), suprathermal electron angle scan
(STEA), and photoelectrons (PHE). Data modes can be run individually or
as part of a data cycle. The data cycles are; normal and calibration.
The normal cycle is a sequence of 15 alternating NSWE/STEA modes
(totaling 30 modes) concluding with a single PHE. The number of
NSWE/STEA mode pairs, and therefore the duration of this cycle, is
command adjustable. Like the ion, the normal cycle runs continuously
once selected. The calibration cycle consists of a sequence of two CALE
modes and works in exactly the same way as the ion calibration cycle.

With the exception of the ENGE data mode, all modes collect data for
four major frames or 64 s. The data collected during this period
comprises one full spectrum subdivided into five 12.8 s spectral
segments. As with the ion modes, this data collection interval ensures
that each spectral segment covers more than a full spin period and
ground analysis software is used to orient the data. A 16 bit to eight
bit count compression algorithm is also implemented as described above.
Because SWEPAM-I and -E are run independently, the data modes are not
synchronized between the two instruments.

**5.3.1 Solar Wind Electrons (NSWE)**

The analyzer plate voltage for the electron sensor has 32 discrete
voltage levels (0 - 31). The first eight levels (0 to 7) correspond to
the photoelectron mode on Ulysses and are not used on ACE. The NSWE data
mode utilizes the highest twenty of these levels (12 through 31),
resulting in electron energy steps centered at 3.2 eV to 1377 eV. During
each 12.8 s spectral segment four E/q levels are sampled, at each of 30
φ-angles, and for each of the seven CEMs. At the nominal 5 rpm spin
rate, angle sectors at a given E/q level are spaced by 12.8° and are
3.0° wide.

The resulting DNSWE data set consists of 20 energy levels at each of 30
φ-angles for the inner most five CEMs (2 - 6) and 15 φ-angles for the
two outermost (largest polar angle) CEMs (1 and 7). The 15 φ-angles of
the outer two CEMs are the result of summing counts from adjacent
φ-angle sectors after the data collection.

**5.3.2 Suprathermal Electron Angle Scan (STEA)**

The STEA data mode utilizes the highest ten voltage levels (22 through
31) corresponding to electron energies centered at 84 through 1377 eV.
During each 12.8 s spectral segment two E/q levels are sampled, at each
of 60 φ-angles, and for each of the seven CEMs. At the nominal 5 rpm
spin rate, angle sectors at a given E/q level are spaced by 6.4° and are
3.0° wide.

The resulting DSTEA data set consists of ten energy levels at each of 60
φ-angles for the inner most five CEMs (2 - 6) and 30 φ-angles for the
outermost two CEMs (1 and 7). The 30 φ-angles of the outer two CEMs are
the result of summing counts from adjacent φ-angle sectors after the
data collection. Comparison of the NSWE and STEA data sets shows how
energy resolution and angular resolution are traded off. In the NSWE
data set, 20 energy levels are measured, but only 30 φ-angles. In the
STEA data set, ten energy levels are measured, but the resolution in φ
is increased by a factor of two. This higher angular resolution is
sometimes necessary for resolving narrow electron beams.

**5.3.3 Photoelectrons (PHE)**

This data mode is identical to the NSWE, with the exception that the
analyzer plate voltages 8 through 27 are used, corresponding to electron
energy steps centered at 0.86 to 429 eV. Thus the PHE mode includes four
low energy steps not used in the NSWE mode. The DPHE data set is
identical in format to the DNSWE data set.

**5.3.4 Electron Calibration (CALE)**

The CALE data mode performs the same function as the CALI mode performs
for SWEPAM-I. In this case the E/q level is fixed at 10 (nominally 25 eV
electrons). The CALE mode data collection consists of 7 CEMs × 5 CEM
bias levels × 51 φ-angles × 2 threshold levels. The discriminator
threshold is alternated between A and B in successive counting intervals
(φ-angles). As with SWEPAM-I, the operating level of the CEM bias supply
(CEMLev) is set by command but varied around this set point by the CALE
data collection routine.

The DCALE data consists of counts from all seven CEMs but unlike
SWEPAM-I, no averaging is required. Thus, the telemetered data set
consists of 7 CEMs × 5 CEM bias levels × 2 threshold levels ×  51
φ-angle counts.

**5.4 Real Time Solar Wind**

The Real Time Solar Wind (RTSW) data set is intended to provide a rough
measure of the solar wind plasma conditions once every 64 s, in real
time, to the international space physics community. To complement this,
magnetic field vector from the MAG instrument, as well as low- energy
particle data from the EPAM instrument and high energy particle data
from the SIS instrument are also being provided in real time. By mutual
agreement of the ACE project and lead investigators for these four
instruments, collection and real-time distribution of these data is
being performed by the Space Environment Center (SEC/NOAA) in Boulder.
It is anticipated that both plots of current solar wind conditions as
well as data sets covering some recent interval of time will be provided
via the World Wide Web (see Zwickl et al., this issue). Together, these
observations should be very useful for real-time monitoring and
prediction purposes and should greatly enhance the nation's ability to
predict and monitor the space weather environment near Earth.

The SWEPAM Real Time Solar Wind data set consists of a subset of the
DSWI data set. These data are only provided when SWEPAM-I is operating
in the SWI mode, and thus will typically provide measurements for 30 out
of every 31 data sets. Of the 96 pixels telemetered as the DSWI data set
in the ACE Science Format , 33 are used for the ACE RTSW format. All
pixels chosen for this purpose have offset angles less then 18°. Since
this represents 34% of the total number of telemetered pixels (or 4.4%
of the total number of measured pixels) any derived parameters (e.g.
density, temperature, speed) must be viewed as incomplete values. While
still extremely valuable for monitoring and prediction, these values
will not be up to full, scientific analysis and publication purposes.
Thus, we advocate the use of a prefix such as 'pseudo-' or 'real-time-'
to delineate these parameters from the full density, temperature, and
speed parameters. Researchers should contact the SWEPAM science team for
complete and properly verified data sets for any scientific purposes.

**Acknowledgments** This paper is dedicated to John Paul Glore who
developed analog electronics at Los Alamos National Laboratory starting
in 1946 and ending with the development of the Ulysses solar wind plasma
experiments. After retiring, JP continued to help out on SWEPAM, which
still contains his amplifier/discriminator circuits, until his death in
1996. The design, development, construction, and testing of space
instrumentation can only be carried out with the combined technical
expertise and dedication of a substantial team effort; we gratefully
acknowledge such critical contributions to SWEPAM from Nancy Baca, Juan
Baldonado, Richard Bramlett, Leroy Cope, Randy Edwards, Danny Everett,
Paul Glore, Irma Gonzales, Jim Lake, Gerrard Martinez, Ignacio Medena,
and Nick Olivas, all at Los Alamos National Laboratory and T.A. Ashlock,
J.C. Chavez, J.W. Daniels, T.J. Ellis, K.M. Olsberg, D.O. Smallwood,
J.H. Temple, and F.J. Wymer at Sandia National Laboratory. This work was
carried out under the auspices of the United States Department of Energy
with financial support from the NASA, Advanced Composition Explorer
program.

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**Figure Captions**

Figure 1. Polar plot of solar wind speed measured by the SWOOPS
experiment on Ulysses over nearly one complete polar orbit about the
Sun. Cyclic variations between high and low speeds, observed on the
right side of the plot at mid-latitudes, are due to rotation of the Sun.

Figure 2. Heliographic longitudes, latitudes, and radial distances of
ACE and Ulysses, spanning the period from the launch of ACE in late 1997
through 2002.

Figure 3. Identification of the solar wind flow state, as given by the
proton bulk speed and density (projected to 1 AU), measured by Ulysses
during its fast latitude scan, compared with the configuration of the
solar corona, as measured in soft X rays by Yohkoh. Yohkoh data was
kindly provided by the Yohkoh mission of ISAS, Japan.

Figure 4. Evidence for magnetic reconnection within the streamer belt as
shown by SWOOPS observations of double proton and alpha-particle streams
on either side of (top and bottom panels), but not at the center of (the
center panel) the heliospheric current sheet.

Figure 5. Color coded plots of suprathermal electron counts for a CME
(C, D) observed by Ulysses on 29 Nov. 1990. Region B displays a heat
flux dropout, which may signify a region that is magnetically
disconnected from the corona at either end while regions A and E show
typical mono-directional electron heat fluxes. Figure taken from McComas
et al. (1994).

Figure 6. Schematic diagram of a CME that drives a shock wave in the
ambient upstream solar wind. The hypothetical disturbance that is
sketched contains both a plasmoid (magnetically disconnected from the
Sun) and a bottle (magnetically connected to the Sun at two footpoints).
The entire disturbance is filled with isolated patches of He-enhanced
plasma that cross boundaries between the three different magnetic
topologies present in interplanetary space.

Figure 7. Schematic illustrations of four generic types of
interplanetary magnetic topologies. At bottom is a simple open topology,
connected at one end to the Sun and at the other end to the interstellar
medium. Heat flux is expected to be uni-directional away from the Sun as
shown by the arrow. Just above is a completely disconnected plasmoid.
Heat flux is expected to be either absent or bidirectional. Next above
is a magnetic bottle or tongue, connected at both ends to the Sun. The
heat flux here is also expected to be bidirectional. At the top is
pictured the simplest flux-rope topology. It is also connected at both
ends to the Sun but contains one or more loops. Figure from Feldman and
Marsch (1997).

Figure 8. Color-coded plot of solar wind electron data for the
corotating interaction region encountered by Ulysses at \~5 AU in
January 1993. Figure taken from Gosling et al. (1993).

Figure 9. The various ionization rates of interstellar hydrogen and
helium atoms calculated from data measured by Ulysses across two
CME-driven shock wave disturbances. Identification of the ionization
processes for the different curves is given in the legend. Figure taken
from Feldman et al. (1995).

Figure 10. Schematic diagram of the SWEPAM-I electro-optics. The ion
sensor is tilted with respect to the spin axis such that observations
from CEMs 1-5 should interleave with those from CEMs 6-10. The exit
aperture mask that has been added for SWEPAM is indicated in both images
(see text).

Figure 11. Photograph of the SWEPAM-I instrument. The large cylindrical
housing is the sensor head; the rectangular electronics box behind it
houses the high and low voltage power supplies and control electronics.

Figure 12. Block diagram of the SWEPAM-I electronics.

Figure 13. Color coded plot of transmitted counts through channel 4 of
SWEPAM-I as a function of azimuthal angle and ESA voltage (equivalent to
E/q) for a 5 kV proton beam incident at a polar angle of 20°.

Figure 14. Central energy (a), azimuthal angle (b), and polar angular
(c) cuts through the channel 4 response function of SWEPAM-I. The solid
lines show the calibration data while the dotted curves display
theoretical fits.

Figure 15. Polar angular cut through all 16 channels of SWEPAM-I at 0°
azimuth and the peak of the E/q transmission curve. The four unmasked,
suprathermal channels have appreciably larger transmissions than they
otherwise would for such large polar angles.

Figure 16. Schematic diagram of the SWEPAM-E electro-optics, similar to
Figure 10.

Figure 17. Photograph of SWEPAM-E; similar to Figure 11.

Figure 18. Block diagram of the SWEPAM-E electronics.

Figure 19. Similar to Figure 13, but for channel 5 of SWEPAM-E,
calibrated with a 1.05 kV beam and an incident polar angle of -21°.

Figure 20. Central energy (a), azimuthal angle (b), and polar angular
(c) cuts through the channel 5 response function of SWEPAM-E along with
theoretical fits, similar to Figure 14.

Figure 21. Polar angular cut through all seven channels of SWEPAM-E at
0° azimuth and the peak of the E/q transmission curve, similar to Figure
15.

Figure 22. Comparison of SWOOPS & SWEPAM-E geometric factors. The
differences in channels 3 and 4 are probably traceable to slight
differences in the actual hardware and CEMs.

Figure 23. Analyzer plate voltage sequence for SWI (track) mode. Azimuth
angle (or spin) is plotted along the x-axis and energy is plotted along
the y-axis. Forty levels are stepped through (5 spectral segments × 8
steps during each φ-angle). This subset of 200 possible levels is
controlled by continually tracking the proton peak. The stepping
sequence for the SSTI (search) mode is qualitatively the same.

Table 1. ACE Level 1 Science Objectives

CRIS ULEIS SWIMS EPAM MAG

SIS SEPICA SWICS SWEPAM

Composition of Matter

Generate table of solar isotropic abundance P P P P C C C

Determine coronal composition P P P P P C P

Compare cosmic ray and solar isotope pattern P P P P P C C C

Measure interstellar/interplanetary pickup ions C C P C C

Determine anomalous cosmic ray composition P P P C C

Origin/Evolution of Elements

Identify solar/metoritic composition differences P P P P P C C

Solar particle contributions to moon/planets/meteorites P P C P P C C

Identify cosmic ray nucleosynthesis process P P

Determine age of cosmic ray source material P P

Search for evidence of galactic evolution P P P C P C

Corona Formation/Solar Wind Acceleration

Isolate coronal formation processes P P P P P C P

Study solar plasma conditions P P P P C

Study solar wind acceleration/fractionation P P P P

Particle Acceleration/Transport

Fractionation in solar flare/interplanetary acceleration P P P P P C P C

Constrain particle acceleration models P P P P P P P P P

Test 3 He-rich and gamma ray flare models C P P P P P C

Measure cosmic ray acceleration/transport time scale P P

Test anomalous cosmic ray origin P P P P

P = PRIMARY MEASUREMENTS

C = CONTRIBUTING MEASUREMENTS

Acronym Name Lead Institution

CRIS Cosmic Ray Isotope Spectrometer CIT

EPAM Electron, Proton, and Alpha-Particle monitor APL

MAG Magnetic Field Monitor Bartol

SEPICA Solar Energetic Particle Ionic Charge State Analyzer UNH

SIS Solar Isotope Spectrometer CIT

SWEPAM Solar Wind Electron, Proton, and Alpha-Particle Monitor Los
Alamos

SWICS Solar Wind Ion Composition Spectrometer UMD

SWIMS Solar Wind Ion Mass Spectrometer UMD

ULEIS Ultra Low Energy Ion Spectrometer UMD

Table 2. SWEPAM Instrument Capabilities

SWEPAM-I SWEPAM-E

Particle species measured 3-D protons & alphas 3-D electrons

Energy range 260 eV/q - 36 keV/q 1.6 eV - 1,350 eV

Energy resolution ∆E/E (FWHM) 5% or 2.5% (1) 12%

Analyzer K (ion energy/ESA voltage) 17.1 4.3

Polar angle FOV (-)25˚ to 65˚ 10˚ - 170˚

Polar angle resolution (FWHM) 5˚ or 2.5° (2) 21˚

Azimuthal angle resolution (FWHM) 3˚ to 4.5˚ (3) 9˚ to 28˚ (3)

G-Factor/pixel (cm^2^ sr eV/eV) 1-20 x 10^-6^ 2-7 x 10^-4^

Time resolution 64 s 64 s

\(1\) 2.5% for energy interleaved data (128 s resolution)

\(2\) 2.5° if orientation of S/C spin axis allows interleaving data from
both sides of fan

\(3\) smaller angles are for normal incidence; larger angles are at the
maximum polar ranges

\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\--

Table 3. SWEPAM Hardware Parameters

SWEPAM-I SWEPAM-E

Box Size (LxWxH) \[cm\] 36 x 24 x 30 25 x 18 x 19

Mass \[kg\] 3.7 2.5

Power, average \[W\] 3.1 2.7

Power, peak \[W\] 3.3 2.9

Telemetry rate \[b/s\] 540 460

Number of CEMs 16 7

[Temperature limits \[°C\]]{.underline}

Preferred operating 0 to +20 0 to +20

In calibration -20 to +45 -20 to +45

Operating survival -25 to +50 -25 to +50

Non-operating survival -30 to +60 -30 to +60

[EMC interference]{.underline}

DC magnetic .01 nT @ 10' .03 nT @ 10'

AC magnetic BDL BDL

AC electrical BDL BDL

Ordinance 2 dimple motors 2 dimple motors

Red tags (HV safe/arm) 2 2

BDL: Below Detectable Levels

Table 4. Geometric factors for SWEPAM-I \[10^-6^ cm^2^ sr eV/eV\].

CEM G CEM G CEM G CEM G

1 7.62 5 8.50 9 0.92 13\* 16.67

2 7.42 6 6.01 10 1.59 14\* 15.20

3 9.23 7 5.07 11 2.86 15\* 20.77

4 8.67 8 1.63 12 4.69 16\* 14.88

\* used for measuring suprathermal instead of solar wind ions

\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\--

Table 5. Geometric factors for SWEPAM-E \[10^-6^ cm^2^ sr eV/eV\].

CEM G CEM G

1 255.6 5 733.5

2 511.0 6 540.3

3 633.4 7 272.9

4 659.5

\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\-\--

Table 6. SWEPAM science data products

+----+---+----+------------------+------------+-----------------------+
|    |   |    |                  |            | Time Res./            |
+----+---+----+------------------+------------+-----------------------+
| A  | D | S  | Description      | Data       | Cadence (s)           |
| cr | a | pe |                  | Matrix     |                       |
| on | t | ci |                  |            |                       |
| ym | a | es |                  |            |                       |
|    | m |    |                  |            |                       |
|    | o |    |                  |            |                       |
|    | d |    |                  |            |                       |
|    | e |    |                  |            |                       |
+----+---+----+------------------+------------+-----------------------+
| DS | S | i  | Plasma Ion       | 40 (E) ×   | 64/64                 |
| WI | W | on | Velocity         | 96 (θ-φ)   |                       |
|    | I |    | Distribution     |            |                       |
+----+---+----+------------------+------------+-----------------------+
| DS | S | i  | Suprathermal Ion | 20 (E) × 4 | 64/64                 |
| TI | W | on | Velocity         | (θ) × 6    |                       |
|    | I |    | Distribution     | (φ)        |                       |
+----+---+----+------------------+------------+-----------------------+
| D  | S | i  | Plasma Ion       | 40 (E) ×   | 64/1984               |
| SS | S | on | Velocity         | 96 (θ-φ)   |                       |
| TI | T |    | Distribution     |            |                       |
|    | I |    |                  |            |                       |
+----+---+----+------------------+------------+-----------------------+
| D  | S | i  | Suprathermal Ion | 20 (E) × 4 | 64/1984               |
| ST | S | on | Velocity         | (θ) × 6    |                       |
| I2 | T |    | Distribution     | (φ)        |                       |
|    | I |    |                  |            |                       |
+----+---+----+------------------+------------+-----------------------+
| D  | S | i  | Real Time Solar  | 40 (E) ×   | 64/64                 |
| RT | W | on | Wind             | 33 (θ-φ)   |                       |
| SW | I |    |                  |            |                       |
+----+---+----+------------------+------------+-----------------------+
| D  | N | el | Plasma Electron  | 20 (E) × 5 | 64/128                |
| NS | S | ec | Velocity         | (θ) × 30   |                       |
| WE | W | tr | Distribution     | (φ)        |                       |
|    | E | on |                  |            |                       |
|    |   |    |                  | 2 (θ) × 15 |                       |
|    |   |    |                  | (φ)\*      |                       |
+----+---+----+------------------+------------+-----------------------+
| D  | S | el | Suprathermal     | 10 (E) × 5 | 64/128                |
| ST | T | ec | Electron         | (θ) × 60   |                       |
| EA | E | tr | Distribution     | (φ)        |                       |
|    | A | on |                  |            |                       |
|    |   |    |                  | 2 (θ) × 30 |                       |
|    |   |    |                  | (φ)\*      |                       |
+----+---+----+------------------+------------+-----------------------+
| DP | P | el | Photoelectron    | 10 (E) × 5 | 64/1344               |
| HE | H | ec | Distribution     | (θ) × 30   |                       |
|    | E | tr |                  | (φ)        |                       |
|    |   | on |                  |            |                       |
|    |   |    |                  | 2 (θ) × 15 |                       |
|    |   |    |                  | (φ)\*      |                       |
+----+---+----+------------------+------------+-----------------------+

\*First set of angles refers to CEMs 2 through 6. Second set for CEMs 1
and 7.