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Richardson and Schwadron: The Limits of Our Solar System
443
The Limits of Our Solar System
John D. Richardson
Massachusetts Institute of Technology
Nathan A. Schwadron
Boston University
The heliosphere is the bubble the Sun carves out of the interstellar medium. The solar wind
moves outward at supersonic speeds and carries with it the Sun’s magnetic field. The interstellar
medium is also moving at supersonic speeds and carries with it the interstellar magnetic field.
The boundary between the solar wind and the interstellar medium is called the heliopause. Before
these plasmas interact at the heliopause, they both go through shocks that cause the flows to
become subsonic and change direction. Voyager 1, which recently crossed the termination shock
of the solar wind, is the first spacecraft to observe the region of shocked solar wind called the
heliosheath. The interstellar neutrals are not affected by the magnetic fields and penetrate deep
into the heliosphere, where they can be observed both directly and indirectly. This chapter
discusses these outer boundaries of the solar system and the interactions between the solar wind
and the interstellar medium, and describes the current state of knowledge and future missions
that may answer the many questions that remain.
1.
INTRODUCTION
The limits of the solar system can be defined in many
ways. The farthest object orbiting the Sun could be one. The
farthest planet could be another. We define this limit as the
boundary between the Sun’s solar wind and the interstellar
medium, the material between the stars in our galaxy. The
interstellar medium is not uniform, but is composed of many
clouds with different densities, temperatures, magnetic field
strengths and directions, and flow speeds and directions.
The Sun is now embedded in a region of the interstellar medium called the local interstellar cloud. The solar wind flows
outward from the Sun and creates a bubble in the interstellar medium.
The solar wind pushes into the local interstellar cloud
until it too weak to push any farther, then turns and moves
downstream in the direction of the local interstellar cloud
flow. The plasmas in these two winds cannot mix, since they
are confined to magnetic field lines and magnetic fields cannot move through one another. Neutrals, however, are not
bound by magnetic fields and can cross between regions. So
local interstellar cloud neutrals can and do penetrate deep
into the solar system. The boundary between the solar and
interstellar winds is called the heliopause.
Both the solar wind and the interstellar wind are supersonic; thus a shock forms upstream of the heliopause in each
flow. At these shocks, the plasmas in each wind are compressed, heated, become subsonic, and change flow direction so that the plasma can move around the heliopause (for
interstellar medium) or down the heliotail (for the solar
wind). The shock in the solar wind is called the termination
shock. The shock in the local interstellar cloud is called the
bow shock. The region between the termination shock and
the heliopause is the inner heliosheath and the region between the bow shock and the heliopause is the outer heliosheath.
Plate 11 shows the equatorial plane from a model simulation of the heliospheric system (Müller et al., 2006). The
color coding in the top panel gives the plasma temperature
and in the bottom panel shows the neutral H density. The
plasma flow lines are shown in the top panel. The trajectories of Voyager 1 (V1) and Voyager 2 (V2) are also shown;
V1 is headed nearer the heliospheric nose than V2. We note
that it is only by serendipity that the Voyager spacecraft are
headed toward the nose of the heliosphere; the Voyager trajectories were driven solely by the locations of the planets
that were part of Voyager’s grand tour of the solar system.
The main heliospheric boundaries, termination shock, heliopause, and bow shock are labeled in the upper panel. The
local interstellar cloud is only barely supersonic, so the heating that occurs at the bow shock is small; at this boundary
the flow of local interstellar cloud plasma begins to divert
around the heliosphere. In the solar wind the flow is radially outward until the termination shock is reached. At the
termination shock, the plasma is strongly heated and the
flow lines bend toward the tail. At the heliopause, the local
interstellar cloud and solar wind flow are parallel; plasma
generally cannot cross this boundary. The neutral density
increases at the nose of the bow shock, forming an enhanced
density region known as the hydrogen wall, and decreases
at the heliopause.
We are just beginning to learn about these outer regions
of our heliosphere and how they interact with the local interstellar cloud. Until V1 crossed the termination shock in
2004 at 94 astronomical units (AU) (1 AU is the distance
from the Sun to Earth), we could only guess the boundary
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The Solar System Beyond Neptune
locations. Now we have observed one termination shock
crossing and have several years of heliosheath data, but we
still have much to learn about the other heliosphere boundaries and the global aspects of these boundaries. This chapter
discusses what we know about the interaction of the heliosphere with the local interstellar cloud and these boundary
regions, the unsolved problems, and the future observations
we hope will help answer these questions.
2.
THE HELIOSPHERE
The heliosphere is the region dominated by plasma and
magnetic field from the Sun. No spacecraft has yet crossed
outside this region; estimates of the size of the heliosphere
at the nose (in the upwind direction) range from 105 to
150 AU (Opher et al., 2006). In the tail, the heliosphere
extends many hundreds and probably thousands of AU. The
location of the nose of the heliopause is determined by pressure balance between the outward dynamic pressure of the
solar wind and the pressure of the local interstellar cloud,
which has contributions from the dynamic pressure, thermal
pressure, and magnetic pressure (see Belcher et al., 1993).
The dynamic pressures are given by 1/2ρV2 where ρ is the
mass density and V is the speed. For the local interstellar
cloud the magnetic pressure B2/8π and the thermal pressures
nkT could be important but the magnitudes are not well
constrained. We discuss first the unperturbed flows from the
Sun and in the local interstellar cloud, then discuss the interactions of these two flows.
2.1.
The Solar Wind
The Sun is a source of both photons and plasma; the
fluxes of each fall off as 1/R2. The effects of solar radiation
on the heliosphere/local interstellar cloud interaction are
small compared to those of the solar wind. The first detections of the solar wind in the early 1960s (Neugebauer and
Snyder, 1962; Gringauz, 1961) confirmed Parker’s (1958)
theoretical predictions of the basic solar wind structure. The
solar wind has been thoroughly measured and monitored
by spacecraft since that time. In addition to the near-equatorial spacecraft in the inner heliosphere, Pioneers 10 and
11 and V1 and V2 have observed the solar wind in the outer
heliosphere, in the case of V1 past 100 AU. Ulysses has provided measurements in the high-latitude heliosphere. These
observations give a good picture of the solar wind parameters throughout the heliosphere.
The solar wind is a supersonic flow of mostly protons,
with on average 3–4% He++ (this percentage varies over a
solar cycle) (Aellig et al., 2001) and a small amount of
highly charged heavy ions. The Sun has a magnetic field
that reverses polarity every 11-year solar cycle. To first order, plasma cannot cross magnetic field lines, so the solar
wind outflow drags the solar magnetic field lines outward.
This outward motion combined with the Sun’s 27-day rotation causes the magnetic field lines to form spirals; at Earth
the spiral angle is about 45°. In the outer heliosphere these
spirals become tightly wound and the magnetic field angle
is nearly 90° (the field is almost purely tangential). Since
the magnetic field lines from the northern and the southern
hemispheres of the Sun have opposite polarity, a heliospheric current sheet (HCS) forms near the equator, across
which the magnetic field direction reverses. The HCS passes
through the termination shock and has been observed in the
heliosheath past 100 AU. The average solar wind speed at
Earth is about 440 km/s with a density of about 7.5 cm–3,
an average temperature of 96,000 K, and an average magnetic field strength of about 5 nT; these values are based
on a solar cycle of data from the Wind spacecraft from 1994
to 2005. The average dynamic pressure is about 2.25 nP;
this pressure determines how big a cavity the solar wind
carves into the local interstellar cloud.
2.2.
The Local Interstellar Cloud
The local interstellar cloud surrounds the heliosphere. It
contains neutral and ionized low-energy particles, H, He, O,
H+, He+, and O+, as well as very-high-energy galactic cosmic
rays (GCRs). These GCRs are observed at Earth but their
intensities are modulated (reduced) by the solar wind magnetic field as they move through the heliosphere. Plate 11
shows how the low-energy plasma deflects around the heliosphere. Helium is the only major species that makes it to
Earth in pristine form and thus provides the best determination of the properties of the interstellar medium. Observations of the H and He emission combined with models
of the neutral flow give estimates of the neutral speed, density, and temperature as well as the flow direction.
Hydrogen is the largest component of the interstellar medium, but it is difficult to determine the local interstellar
cloud properties from H observations because a large fraction of the H charge exchanges with the local interstellar
cloud H+ before it enters the heliosphere. Charge exchange
occurs when an ion and a neutral collide and an electron
goes from the neutral to the ion. The former ion is a now
neutral and is not bound by magnetic fields, so it continues
to move at the speed it had when it was an ion. The former
neutral, now an ion, is bound to the magnetic field and
accelerated to the speed of the bulk plasma flow. Outside
the bow shock the plasma and neutrals have the same speed
and temperature so charge exchange has no effect. But inside the bow shock, the plasma parameters are different
from those in the interstellar medium since the plasma is
heated, compressed, and deflected at the shock. The plasma
slows further as it piles up in front of the heliosphere. So
when charge exchange occurs, the new neutrals have speeds
and temperatures from the shocked plasma, not the pristine
interstellar medium. The H observed in the heliosphere has
an inflow speed of only 20 km/s, 6 km/s less than the interstellar medium speed (Lallement et al., 1993), confirming
that the H has charge exchanged with local interstellar cloud
H+, which has slowed down near the heliosphere boundary.
The plasma flow in this region has begun to divert around
the heliosphere, so the flow direction of the H is also altered.
Richardson and Schwadron: The Limits of Our Solar System
Observations with the SOHO/SWAN instrument show that
the arrival direction of the interstellar H differs from that
of the interstellar He by about 4° (Lallement et al., 2005).
This offset results from charge exchange of the H with the
deflected H+ in front of the heliosphere. The direction of
the offset tells us the direction of the plasma flow in the
local interstellar cloud just outside the heliopause and thus
provides information on the magnetic field direction (Lallement et al., 2005). The magnetic field strength in the local
interstellar cloud is not known; generally it is thought to
be on the order of 1 µG with an upper limit of 3–4 µG.
The interstellar neutrals enter the heliosphere and absorb
and emit sunlight at specific wavelengths. Backscattered H
Lyman-α (Bertaux and Blamont, 1971; Thomas and Krassa,
1971) and He 58.4 nM radiation have both been observed
(Weller and Meier, 1979). As the neutrals move toward the
Sun, they are accelerated inward by gravity but pushed outward by the radiation pressure (see review by Fahr, 1974).
For H, the radiation pressure and gravity are comparable,
so the H is not accelerated inward. Helium is heavier and
is accelerated inward. Hydrogen also has a much shorter
ionization time in the solar wind than He. The result is that
most H becomes ionized before it reaches 5–10 AU. The
peak in backscattered Lyman-α intensity is roughly in the
direction of the local interstellar cloud flow with a small
deflection occurring from the charge exchange in the hydrogen wall.
The ionization time of He is long; most He is not ionized
until it is inside 1 AU. The He with trajectories that remain
outside 1 AU survive and flow downstream out of the heliosphere. The Sun’s gravity bends the trajectories so that the
He is focused downstream of the Sun, where the maximum
emission from He is observed (Axford, 1972; Fahr and Lay,
1973). The location of the maximum He emission is thus
opposite to the direction of the incoming He flow. Observation of the direction with peak He intensity is one method
used to determine the local interstellar cloud flow direction.
The interstellar neutrals are ionized in the solar wind via
charge exchange with protons (the dominate process for H)
or photoionization. These neutrals can be observed directly
as H+ and He+ pickup ions. He+ is rarely observed in the
solar wind and thus is a good marker of interstellar neutrals.
Both H+ and He+ pickup ions have distinctive signatures in
the solar wind; when the neutral H and He become ionized
they are accelerated to the solar wind speed. These ions,
called pickup ions because they are picked up (accelerated)
by the solar wind, have a gyromotion roughly equal to the
solar wind speed, so these ions show a sharp energy cutoff
at four times the solar wind energy.
The energy to accelerate the pickup ions comes from the
bulk flow of the solar wind, so an indirect way to measure
the pickup ion density (discussed further below) is to look at
how much the solar wind slows down as it moves outward.
These direct and indirect observations can be combined with
models of the neutral inflow and loss to estimate speeds,
temperatures, and densities of H and He in the interstellar
medium.
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The gas detector on Ulysses measures the neutral He
velocity and temperature directly (Witte, 2004; Witte et al.,
1992). Even with these direct measurements, determination
of the density of the local interstellar cloud still requires
modeling since some of the He is lost on its way through the
heliosphere. All these measurements taken together provide
fairly tight constraints on the local interstellar cloud He neutral speed and temperature (Möbius et al., 2004).
The He moves toward the Sun from an ecliptic longitude
of 75° and an ecliptic latitude of –5° with a speed of 26.4 ±
0.3 km/s. The temperature is 6300 ± 340 K and the density
of He is 0.015 ± 0.002 cm–3. The collision timescales in the
local interstellar cloud are less than the neutral and plasma
lifetimes, so the plasma and neutral species should all have
the same speed and temperature far from the heliosphere.
Since the H is altered by charge exchange with the plasma
near the heliosphere, the H density is more difficult to determine. Best estimates are on the order of 0.20 ± 0.03 cm–3
in the local interstellar cloud. Since the ionization fraction
depends on the electron temperature, which is not well
known, the plasma density is also uncertain, with best estimates of 0.06–0.10 cm–3 based on the speed of the H measured at Earth (Lallement et al., 1993).
3. THE TERMINATION SHOCK LOCATION
Although some properties of the local interstellar cloud
are well constrained, others are not, which makes determining the heliosphere size and shape difficult. We have observed one boundary location; the termination shock at the
position of V1 was at 94 AU on December 16, 2004 (Stone
et al., 2005; Decker et al., 2005; Burlaga et al., 2005). This
crossing provides one point with which to calibrate our
models. Since the average solar wind dynamic pressure is
2.25 nP at 1 AU, one can roughly estimate that the total
local interstellar cloud pressure must be 2.5 × 10–4 nP. But
there are problems with this simple approach, which we describe below.
3.1.
Solar Wind Time Dependence
The average solar wind values given above are only
a piece of the story. The solar wind pressure is variable
on scales from hours to 11-year solar cycles and likely has
longer-term variations. Small-scale variations of the pressure do not affect the boundary locations, but larger-scale
variations change the balance between the solar wind and
local interstellar cloud pressures and cause the termination
shock and heliopause to move in and out. The solar wind
dynamic pressure changes by a factor of about 2 over a solar
cycle, with minimum pressure near solar maximum (when
sunspot numbers are highest), then a rise in pressure with
a peak a few years after solar maximum, followed by a decrease to the next solar maximum (Lazarus and McNutt,
1990). Models of the heliosphere show that these pressure
changes drive an inward and outward “breathing” of the heliosphere; the termination shock location changes by about
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The Solar System Beyond Neptune
10 AU over a solar cycle and the heliopause by 3–4 AU
(Karmesin et al., 1995; Wang and Belcher, 1998). The solar
wind structure also changes over the solar cycle. At solar
minimum, the solar wind near the equator is slow (400 km/
s) and dense (8 cm–3). The solar wind at higher latitudes
comes from polar coronal holes and is fast, 700–800 km/s
with a density of 3–4 cm–3. At solar maximum, the polar
coronal holes are not present and most of the flow is lowspeed solar wind. Despite these latitudinal changes in the
solar wind, the dynamic pressure changes over the solar cycle are the same at all heliolatitudes (Richardson and Wang,
1999), so the heliosphere is not expected to change shape
over a solar cycle.
Large events on the Sun can also affect the termination
shock position. Active regions on the Sun generate coronal
mass ejections (CMEs), explosive discharges of material
into the solar wind where they are called interplanetary
CMEs (ICMEs). ICMEs propagate through the heliosphere.
When a series of ICMEs occurs, the later ones catch up
to and merge with previous ICMEs, causing a buildup of
plasma and magnetic field known as a merged interaction
region (MIR) (Burlaga et al., 1984; Burlaga, 1995). These
regions are usually preceded by a fast forward shock at
which the speed, density, temperature, and magnetic field
strength all increase. Shocks are rare in the outer heliosphere; most can be related to a large CME or a series of
CMEs on the Sun (see Richardson et al., 2005a). The dynamic pressure in the MIRs can be a factor of 10 higher
than that in the ambient solar wind and these structures last
1–3 months (Richardson et al., 2003). The MIRs drive the
termination shock outward several AU for 2–3 months, after
which the termination shock rebounds inward (Zank and
Müller, 2003).
3.2.
Local Interstellar Cloud Time Dependence
Changes in the pressure or magnetic field of the local
interstellar cloud could also cause the boundary locations to
move. The scale length for changes in the local interstellar
cloud is not known. The Sun is very near the local interstellar cloud boundary and will cross into another interstellar medium environment in a few thousand years, and scale
lengths for variations may be smaller near this boundary
(Müller et al., 2006). Observations of He by Ulysses are
consistent with no change in the interstellar neutral density
since the launch of Ulysses in 1990 (Witte, 2004).
3.3. Local Interstellar Cloud Effects
on the Solar Wind
If we assume that the local interstellar cloud pressure
were constant, we can use the observed solar wind pressures to model the heliosphere boundary changes. Models
must include the effects of the local interstellar cloud neutrals on the solar wind. These neutrals are the largest (by
mass) component of the interstellar space outside about
Fig. 1. The density of the plasma and neutral components from
1 to 1000 AU. The solar wind ions come from the Sun. The pickup
ions are interstellar neutrals that have been ionized in the solar
wind. The densities of the solar wind and of the pickup ions jump
at the termination shock. Outside the heliopause, the ions are part
of the local interstellar cloud. Both the ion and neutral density
increase in front of the heliopause (in the so-called hydrogen wall).
The interstellar neutrals dominate the mass density outside 10 AU.
Figure courtesy of R. Mewaldt.
10 AU. Figure 1 shows the density of plasma and neutrals
vs. distance (we know now the termination shock and heliopause in this figure are located too close to the Sun). The
solar wind plasma falls off as R–2, then increases abruptly
at the termination shock and heliopause. The H density is
about 0.1 at the termination shock and 0.2 in the local interstellar cloud, then decreases rapidly inside 10 AU, the
ionization cavity inside which most H is ionized. Both the
H and H+ increase between the bow shock and heliopause
and form the hydrogen wall. The pickup ion density/thermal
ion density ratio increases to the termination shock, where
about 20% of the ions are pickup ions. This ratio is consistent with the observed slowdown of the solar wind. Outside
about 30 AU the pickup ions (of local interstellar cloud origin) dominate the plasma thermal pressure. So even though
the Sun has carved out its own space in the local interstellar cloud, the local interstellar cloud influence is very strong.
3.3.1. Solar wind slowdown. The neutral interstellar H
is continuously being lost, mostly by charge exchange with
solar wind protons. The neutrals formed from the solar wind
protons escape the solar system at the solar wind speed. The
new protons formed from ionized local interstellar cloud
neutrals are accelerated from their initial speed of 20 km/s
inward to the solar wind speed, 400–700 km/s outward. As
mentioned earlier, plasma does not flow across field lines;
ions at speeds different from the bulk solar wind speed are
accelerated by an electric field E = –VXB. The initial thermal energy of the pickup ions is equal to the solar wind en-
Richardson and Schwadron: The Limits of Our Solar System
ergy, about 1 keV. The energy for the acceleration and heating comes from the bulk flow of the solar wind.
Two consequences of these pickup ions are that the solar wind is heated and slows down. The amount of the slowdown depends on the density of the local interstellar cloud
H, so measurements of the slowdown give an estimate of the
local interstellar cloud density. This slowdown can be difficult to measure because the solar wind speed varies with
time and solar latitude, as well as with distance. In the inner heliosphere, the slowdown is small and difficult to detect. In the outer heliosphere, determining the slowdown
from the inner heliosphere to V2 requires (1) another spacecraft in the inner heliosphere at a similar heliolatitude to
provide a speed baseline and (2) a model to propagate the
solar wind outward to determine how much of the speed
change comes from radial evolution of the solar wind and
how much is due to pickup ions. Figure 2 shows a compilation of various determinations of the slowdown. The plot
shows the solar wind speed normalized to 1 AU. The horizontal lines show the average amount of slowdown; the
length of these lines shows the radial range covered in each
study. The vertical spread shows the uncertainties in the
slowdown determination. The curves shows the predicted
solar wind slowdown for densities of interstellar H at the
termination shock of 0.07, 0.09, and 0.11 cm–3. The data
are best fit with a density of about 0.09 cm–3 at the termi-
nation shock with an uncertainty of 0.01 cm–3. We note that
the H density at the termination shock is not the density in
the local interstellar cloud; much filtration, or H loss, occurs between the bow shock and the termination shock.
By 95 AU, the solar wind speed has decreased by about
20% [and since the mass flux, the speed (v) times the mass
density (ρ), is a constant, the density must increase by 20%].
So the dynamic pressure, ρv2, is only 77% of what it was at
1 AU. The effect of the interstellar H is to lower the solar
wind dynamic pressure. This reduced pressure causes the
heliosphere boundaries to be closer to the Sun. The solar
wind slowdown is related to the pickup ion density by npu/
nsw = (3γ – 1)/(2γ – 1)(δv/v0), where nsw is the solar wind
density, npu is the pickup ion density, δv is the solar wind
slowdown, v0 is the solar wind speed at 1 AU, and is the
ratio of specific heats. For γ = 5/3, npu/nsw = 7δv/6v0
(Richardson et al., 1995).
3.3.2. Solar wind heating. As the solar wind moves
out, it expands and thus might be expected to cool. The solar
wind temperature does decrease out to about 30 AU, but
not as quickly as adiabatic cooling would predict, then increases with distance with some solar cycle variations superposed. Internally driven processes, such as shocks and
dissipation of magnetic fluctuations, provide some heating,
but not enough to produce the observed temperature profile (Richardson and Smith, 2003; Smith et al., 2001, 2006).
The source of this heating is the pickup ions. These ions
dominate the thermal energy of the plasma outside about
30 AU, and a small fraction of their energy is transferred
to the thermal plasma. The mechanism for this energy transfer is through waves; the pickup ions initially have ring
distributions, which are unstable. Low-frequency waves are
generated that isotropize the pickup ions, but about 4% of
the energy in these waves is transferred to the thermal ions,
heating the solar wind (Smith et al., 2006; Isenberg et al.,
2005). Since the pickup ion energy and thus the amount of
energy transfered to the thermal ions are proportional to the
solar wind speed, the solar wind speed and temperature are
strongly correlated.
3.4.
Fig. 2. The plot shows the slowdown of the solar wind caused
by the interstellar neutrals. The fraction of the initial solar wind
speed is plotted vs. distance. The horizontal lines with error bars
show observed speed decreases; the lengths of these lines show
the radial width over which the slowdown was observed. The
vertical lines are error bars. The curves shows model predictions
of the solar wind speed as a function of distance for interstellar
neutral H densities of 0.09 cm–3 (solid line), 0.07 cm–3 (upper
dashed line), and 0.11 cm–3 (lower dashed line).
447
Termination Shock Model Results
Many models developed to determine the termination
shock motion, the asymmetry of the boundaries, and the distance between boundaries (Webber, 2005; Linde et al., 1998,
Pogorelov et al., 2004; Izmodenov et al., 2005; Opher et al.,
2006). We have developed a two-dimensional model that
includes the effect of the local interstellar cloud neutrals on
the solar wind and use it to track the termination shock location. Figure 3 shows predictions of the termination shock
location based on a two-dimensional model that includes the
effect of the local interstellar cloud neutrals on the solar
wind. The model input is the solar wind pressure observed
by V2; the profile is normalized to the V1 termination shock
crossing (Richardson et al., 2006a). The model shows that
V1 and the shock were both moving radially outward start-
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The Solar System Beyond Neptune
Fig. 3. The predicted position of the termination shock using V2
observations of the solar wind as input to a two-dimensional timedependent hydrodynamic model. The termination shock locations
are extended into the future by repeating the pressure profile from
the previous solar cycle. The V1 and V2 trajectories are shown
by the dashed diagonal lines. The lower curve shows the termination shock position shifted inward 10 AU to simulate a possible
termination shock asymmetry. The termination shock moved inward from 1995 to 1999, then moved outward ahead of V1 until
2004, when it moved inward and crossed V1. The model suggests
V2 could cross the termination shock before the end of 2007.
ing in 1999, that V1 made some close encounters to the
shock, but that the shock crossing did not occur until the
termination shock moved inward in 2004.
3.5.
Heliospheric Asymmetries
At Earth, the 45° average angle of the solar magnetic
field to the solar wind flow direction results in asymmetries
in the magnetosheath plasma (Paularena et al., 2001) and
the magnetopause locations (Dmitriev et al., 2004). The
draping of magnetic field lines around the magnetosphere
results in different field strengths on the dawn and dusk
sides, causing these asymmetries. The interstellar magnetic
field could result in a similar distortions of the termination
shock and heliopause and significantly affect the width of
the heliosheath (see, e.g., Linde et al., 1998; Ratkiewicz et
al., 1998; Pogorelov et al., 2004). Voyager 2 entered the termination shock foreshock region much closer to the Sun
(75 AU) than V1 (85 AU) (Stone et al., 2005). This large a
difference probably does not arise entirely from solar wind
dynamic pressure changes but also from asymmetries in the
boundary locations, with the boundaries closer in the direction of V2.
The nature of the heliosphere distortion depends on the
orientation and strength of the interstellar magnetic field.
Several suggestions for the magnetic field direction in the
local interstellar cloud near the heliosphere have been published. The location of the sources of the heliospheric radio
emissions (described below) lies along a line roughly parallel to the galactic plane. Gurnett et al. (2006) suggest that
the radio emission should occur on magnetic field lines tan-
gential to the shock that excites these emissions. In this case,
the magnetic field would be roughly perpendicular to the
galactic plane. A field orientation 60° from the galactic
plane is derived from differences in the flow velocities of
interstellar H and He (Lallement et al., 2005; Izmodenov et
al., 2005).
Magnetohydrodynamic (MHD) models of the heliosphere that include the effects of an interstellar magnetic
field reveal that the distortion of the heliosphere is sensitive
to the angle between the field and the velocity of the interstellar wind. Opher et al. (2007) show that a field tilted 60°–
90° from the galactic plane is needed to match the flow data
and the radio emission data. This tilt creates a large asymmetry in the model heliosphere in both the north/south and
east/west directions. The models show that magnetic field
lines that intersect the termination shock extend 2 AU inside the termination shock at the position of V1 and 4–7 AU
inside the termination shock in the direction of V2 (Opher et
al., 2006). These distances represent the average thickness
of the foreshock region in front of the termination shock.
Small-scale magnetic field variations will affect these distances locally. This model also predicts that the distance of
the termination shock in the direction of V2 is 8–11 AU
closer than in the direction of V1 and the heliopause distance is 13–24 AU closer in the direction of V2. This model
did not include neutrals, which may reduce the asymmetry
(Pogorolev et al., 2006). Other models also give asymmetries with ratios of the termination shock distances at V1
and V2 varying from 1.05 to 1.33. The predicted heliopause
locations vary from 145–157 AU at V1 and 109–139 AU at
V2, with heliosheath thicknesses of 55–67 AU toward V1
and 33–50 AU toward V2 (see Opher et al., 2006).
None of these models include all the important physics;
that task is currently too expensive computationally. In particular, the effects of the HCS may be important, since the
location of the HCS at the termination shock varies over a
solar rotation, but this small-scale structure is hard to include in a global model. Although the models differ quantitatively, the results suggest that the average shock distance
at V2 could be ~5–10 AU closer than at V1. In addition,
the solar cycle variation of the solar wind dynamic pressure
is such that in 2007 the shock should be several AU closer
than average, suggesting that V2 might cross the shock in
2007–2008. The models also indicate that the heliosheath
is likely narrower at V2, so V2 may cross the heliopause at
about the same time as V1. Voyager 1 and V2 observations
of the asymmetry of the termination shock and heliosheath
should make possible improved models and estimates of the
distances to the heliopause.
4. COSMIC RAYS
Cosmic rays are high-energy charged particles that bombard Earth from above the atmosphere. These particles have
important effects at Earth. Cloud cover is directly correlated
with cosmic ray flux, so these particles may affect Earth’s
weather (Svensmark et al., 2007). Several thousand cosmic
Richardson and Schwadron: The Limits of Our Solar System
rays pass through a person’s body every minute. These can
cause biologic damage but also cause mutations that accelerate evolution. A new source of cosmic rays discussed below is generated from the grains in the Kuiper belt (Schwadron, 2002). These cosmic rays were likely more intense in
the past and the termination shock was at times much closer
to the Sun than it is today. Both of these effects may have
caused periods when Earth was ensconced in much more
severe radiation environments than at present. Several types
of cosmic rays come from the outer heliosphere, as discussed below.
4.1.
Galactic Cosmic Rays
Galactic cosmic rays (GCRs) possess energies above
100 MeV and are thought to be accelerated in the shock
waves associated with supernova in our galaxy. The GCRs
are predominately fully stripped nuclei but lesser fluxes
of GCR electrons are also observed. Galactic cosmic rays
representing most elements in the periodic table have been
detected. Since cosmic rays are charged, they are affected
by magnetic fields in interplanetary space, the heliosphere,
and near Earth. Although GCRs do reach Earth, many are
deflected by the heliospheric magnetic field before they
reach the inner solar system. Fewer GCRs are observed at
Earth at solar maximum, when the solar wind magnetic field
is stronger and more turbulent. Large merged interaction
regions (MIRs) temporarily reduce the flux of GCRs as
these structures contain large and turbulent magnetic fields.
An empirical relation between GCRs and the magnitude of
B shows that the GCR flux increases when B is above normal and increases when B is below normal; this relationship
holds out to the termination shock (Burlaga et al., 2003a,
2005). Spacecraft in the outer solar system see increased
GCR fluxes. These particles move so fast and are scattered
so thoroughly by the interstellar magnetic field that their
flux is likely to be fairly uniform outside the heliosphere.
4.2.
Anomalous Cosmic Rays
Early cosmic-ray observers discovered an unusual subset
of cosmic rays that consisted of singly ionized ions (instead
of fully stripped nuclei) with energies of 1–50 MeV/nuc (see
Mewaldt et al., 1994). Most of the anomalous cosmic rays
(ACRs) are species that have high ionization thresholds,
such as He, N, O, Ne, and Ar. Until recently, ACRs were
thought to arise only from neutral atoms in the interstellar
medium (Fisk et al., 1974) that drift freely into the heliosphere through a process that has four essential steps: first,
the neutral particles from the local interstellar cloud enter the
heliosphere; second, the neutrals are converted into pickup
ions; third, the pickup ions are preaccelerated by shocks and
waves in the solar wind [see also section 5 and Schwadron
et al. (1998)]; and finally, they are accelerated to their final
energies at or beyond the termination shock (Pesses et al.,
1981). Easily ionized elements such as C, Si, and Fe are
expected to be strongly depleted in ACRs, since these ele-
449
ments are not neutral in the interstellar medium and therefore cannot drift into the heliosphere.
Instruments such as the Solar Wind Ion Composition
Spectrometer (SWICS) on Ulysses are able to detect pickup
ions directly. Ongoing measurements by cosmic-ray detectors on Voyager and Wind have detected additional ACR
components (Reames, 1999; Mazur et al., 2000; Cummings
et al., 2002). We now understand that, in addition to the
traditional interstellar source, grains produce pickup ions
throughout the heliosphere: Grains near the Sun produce
an “inner source” of pickup ions, and grains from the Kuiper belt provide an “outer” source of pickup ions and anomalous cosmic rays.
Recent observations call into question not only the
sources of ACRs, but also the means by which they are accelerated. The prevailing theory, until V1 crossed the termination shock, was that pickup ions were energized at the
termination shock to the 10–100-MeV energies observed
(Pesses et al., 1981). When V1 crossed the termination
shock, it did not see a peak in the ACR intensity as this theory predicts (Stone et al., 2005). Instead, the ACR intensities continued to increase in the heliosheath. One suggestion
is that the ACRs are accelerated in the flanks of the heliosheath and the source region is not observed by V1 near the
nose (McComas and Schwadron, 2006). Another suggestion is that the acceleration region was affected by a series
of MIRs before the V1 crossing that diminished the acceleration process (McDonald et al., 2006; Florinski and Zank,
2006), in which case V2 may see a different ACR profile
when it crosses the termination shock.
4.3.
The Interstellar and Inner Source Pickup Ions
The telltale signature of interstellar pickup ions, as seen
for H+ in Fig. 4, is a cutoff in the distribution at ion speeds
twice that of the solar wind in the spacecraft reference
frame. This cutoff is produced since interstellar ions are initially nearly stationary in the spacecraft reference frame and
subsequently change direction, but not energy, in this frame
due to wave-particle interactions and gyration, causing them
to be distributed over a range of speeds between zero and
twice the solar wind speed in the spacecraft frame.
In comparison to the interstellar pickup ion distributions,
the C+ and O+ distributions in Fig. 4 are puzzling. The distributions do not cut off at twice the solar wind speed and
have a peak, contrary to the other ACR ions, which have
flat distribution. These ions are singly charged, which precludes them from having been emitted directly by the Sun;
if this were the case, they would be much more highly
charged, e.g., C5+ and O6+ are each common solar wind
heavy ions. The C+ and O+ distributions in Fig. 4 are consistent with a pickup ion source close to the Sun. As the ions
travel out in the solar wind, they cool adiabatically in the
solar wind reference frame, and instead of having a distribution that cuts off at twice the solar wind speed in the
spacecraft reference frame, they have a peak near the solar
wind speed. The solid curve that goes through the C+ and
450
The Solar System Beyond Neptune
grains divided by the net area of the sky over which they are
distributed; so, for example, a large solid object would have
a filling factor of 1, whereas a finite set of points would have
a filling factor of 0). A large filling factor is consistent with
observations of the inner source that require a net geometric cross-section typically 100 times larger than that inferred
from zodiacal light observations (Schwadron et al., 2000).
The very small grains act effectively as ultrathin foils for
neutralizing solar wind (Funsten et al., 1993), but are not
effective for scattering light owing to their very small size.
The implications of this suggestion are striking: The inner
source may come from a large population of very small
grains near the Sun generated through catastrophic collisions of larger interplanetary grains (Wimmer-Schweingruber
and Bochsler, 2003).
Fig. 4. Observed distributions from Ulysses/SWICS of C+ (solid
triangles), O+ (open circles), and H+ (open triangles) vs. ion speed
in the spacecraft frame normalized by the solar wind speed
(Schwadron et al., 2000). The observations are compared with
simulated distributions of solar wind H+ (dashed black curve),
interstellar H+ (upper gray dash-dot curve), inner source H+ (upper thick black line), inner source C+, O+ (lower thick black line),
and interstellar O+ (lower gray dash-dot curve).
O+ data points is based on a transport model (Schwadron,
1998; Schwadron et al., 2000) for ions picked up close to
the Sun, thereby proving the source is in the inner solar
system. Data like those presented in Fig. 4 suggested that
another pickup ion source was present in addition to interstellar neutrals (Geiss et al., 1995; Gloeckler and Geiss,
1998; 2000a; Schwadron et al., 2000). Since the source is
peaked near the Sun, as is the distribution of interplanetary
grains, it was hypothesized that grains were associated with
the source.
The initial, naive expectation was that the inner source
composition would resemble that of the grains, i.e., enhanced in C and O with strong depletions of noble elements
such as Ne. However, the composition strongly resembled
that of the solar wind (Gloeckler et al., 2000a). This conundrum was resolved by assuming a production mechanism
whereby solar wind ions become embedded within grains
and subsequently reemitted as neutrals. Remarkably, the
existence of the inner source was hypothesized many years
prior to its discovery (Banks, 1971).
The concept that the inner source is generated due to
neutralization of solar wind requires that sputtered atoms
do not strongly contribute to the source. However, for grains
larger than a micrometer, sputtering yields of the grains
are much larger than the yields of neutralized solar wind
(Wimmer-Schweingruber and Bochsler, 2003). This suggests
that the grains that give rise to the inner source are extremely small (hundreds of angstroms). A small grain population generated through catastrophic collisions of larger
interplanetary grains would also yield a very large filling
factor (the filling factor is the net area of the sky filled in by
4.4.
Pickup Ions from Cometary Tails
Comets are also a source of pickup ions in the heliosphere. Until recently, it was thought that the only way to
observe cometary pickup ions was to send spacecraft to
sample cometary matter directly. However, Ulysses has now
had three unplanned crossings of distant cometary tails,
Comets McNaught-Hartley (C/1991 T1), Hyakutake (C/
1996 B2), and McNaught (C/2006 P1) (Gloeckler et al.,
2000b; Gloeckler, 2004; Neugebauer et al., 2007). These
three unplanned crossings of cometary tails by Ulysses suggest an entirely new way to observe comets through detection of distant cometary tails.
4.5. Outer Source Pickup Ions and
Anomalous Cosmic Rays
Recent observations from the Voyager and Wind spacecraft resolved ACR components comprised of easily ionized elements [such as Si, C, Mg, S, and Fe (Reames, 1999;
Mazur et al., 2000; Cummings et al., 2002)]. An interstellar
source for these “additional” ACRs, other than a possible
interstellar contribution to C, is not possible (Cummings et
al., 2002). Thus, the source for these ACRs must reside
within the heliosphere. A number of potential ACR sources
are present within the heliosphere (Schwadron et al., 2002).
The solar wind particles are easily ruled out as a source
since they are highly ionized, whereas ACRs are predominantly singly charged. Discrete sources such as planets are
also easily ruled out since their source rate is not sufficient
to generate the needed amounts of pickup ions. Another potential source is comets, but the net cometary source rate is
sufficiently large only inside 1.5 AU, a location so close to
the Sun that the generated pickup ions would be strongly
cooled by the time they reach the termination shock, making
acceleration to ACR energies extremely difficult. Moreover,
a cometary source would naturally be rich in C, which is
not consistent with the compositional observations of easily
ionized ACRs.
The inner source (discussed in the preceding section) is
another possibility, but again, adiabatic cooling poses a sig-
Richardson and Schwadron: The Limits of Our Solar System
451
Fig. 5. An illustration of ACR production (Schwadron et al., 2002). Lower curves apply to the known interstellar source ACRs (adapted
from Jokipii and MacDonald, 1995), while upper curves apply to the outer source, described later in the paper.
nificant problem since these particles are picked up so close
to the Sun. It has been suggested that the inner source may
be substantially preaccelerated in the inner heliosphere,
thereby overcoming the effects of adiabatic cooling. Contrary to this suggestion, inner source observations in slow
solar wind show pronounced effects of adiabatic cooling
(Schwadron et al., 1999) and charge-state observations of
energetic particles near 1 AU indicate no evidence for acceleration of the inner source (Mazur et al., 2002). The additional population of ACRs requires a large source of pickup
ions inside the heliosphere and outside 1 AU.
Schwadron et al. (2002) suggest that the source is material extracted from the Kuiper belt through a series of
processes shown schematically by the lines in Fig. 5: First,
micrometer-sized grains are produced due to collisions of
objects within the Kuiper belt; grains spiral in toward the
Sun due to the Poynting-Robertson effect; neutral atoms are
produced by sputtering and are converted into pickup ions
when they become ionized; the pickup ions are transported
by the solar wind to the termination shock and, as they are
convected, are preaccelerated due to interaction with shocks
and due to wave-particle interactions; and finally, they are
injected into an acceleration process at the termination shock
to achieve ACR energies. The predicted abundances are all
within a factor of 2 of observed values, providing strong
validation of this scenario. See Schwadron et al. (2002) for
a detailed presentation of these points.
5.
THE TERMINATION SHOCK FORESHOCK
When V1 reached 84 AU in 2002 it observed increases in
energetic particles with energies from 40 keV to >50 MeV
(Stone et al., 2005; Decker et al., 2005). Figure 6 shows that
these ion fluxes increased in mid-2002 by a factor of 100,
disappeared in early 2003, came back at high energies in
mid-2003 and at low energies in early 2004, and persisted
until the termination shock crossing in late 2004. These
intensity increases were not associated with solar events and
were not observed by V2, which trails V1 by 18.5 AU or
about six years. Thus it seemed likely that these particles
were associated with the termination shock. But these particles were flowing along the magnetic field lines predomi-
452
The Solar System Beyond Neptune
Fig. 6. V1 energetic particle data during 2002.0–2005.7 (83.4–97.0 AU). (a), (b), (d) Background-corrected, scan-averaged five-day
smoothed daily-averaged intensities of 40–53-keV ions, 3.4–17.6-MeV protons, and 0.35–1.5-MeV electrons. (c) First-order anisotropy vector A1/A0 of proton channel in (b) when the intensity >0.0025 flux units (inset shows orientation of R and T). Whiskers
show direction that particles are traveling. (e) Black bars show times when the V1 plasma wave instrument detected electron plasma
oscillations (Gurnett and Kurth, 2005). (f) Estimated plasma radial flow speed VR based on 40–220-keV ion angular data and a velocity extraction algorithm (Krimigis et al., 2005). Bounds on VR during period A and the second half of period C (2004.56–2004.78) are
indicated by the cross-hatched rectangles and means are shown by the horizontal bars. Horizontal bars during 2002.041–2002.172 and
2003.257–2003.322 are from Krimigis et al. (2003). Figure courtesy of R. Decker.
nately in one direction and that direction was away from
the Sun (the actual field direction is tangential due to the
winding of the Parker spiral with distance, so the particles
were moving tangentially but from the direction that is connected to the Sun). If the particle source were the termination shock and V1 was inside this shock, then the particles
seemed to be flowing the wrong way. This backward flow
direction led to suggestions that V1 had already crossed the
termination shock (Krimigis et al., 2003) and that the particles were coming outward from the termination shock.
Subsequent work has shown that if the termination shock
were blunt, the magnetic field lines first intersect the termination shock closer to the Sun at the distance of V1, so that
particles from the termination shock would flow to V1 in
Richardson and Schwadron: The Limits of Our Solar System
the nominally outward direction (Jokipii et al., 2004). This
hypothesis predicts that flows observed by V2, on the opposite side of the nose of the termination shock, would be
in the opposite direction (nominally toward the Sun). In late
2004, V2 observed these particles and they were in the opposite direction from those at V1, as predicted (Decker et
al., 2006). Model predictions (Opher et al., 2006) show that
field lines from the termination shock should intersect V1
2–3 AU upstream of the termination shock and V2 4–7 AU
upstream of the shock. The amount of time each Voyager
spends in the foreshock region depends on the motion of the
termination shock. As V1 approached the termination shock
the solar wind dynamic pressure was increasing, pushing
the termination shock outward, so that V1 was surfing behind the shock and remained in or near the foreshock for
almost 9 AU. The solar wind pressure then decreased and
the shock moved quickly inward past V1.
Figure 6 shows that the first termination shock particle
event ended at 2003.1. The end of this event coincided with
the passage of a merged interaction region (MIR) by the
spacecraft that pushed the termination shock and thus the
foreshock region outward. The first termination shock particle event at V2 also ended when a shock followed by a MIR
passed V2, driving the shock outward. The foreshock particle fluxes observed have a lot of structure and many peaks.
Some of this structure is connected with the arrival of MIRs
at the location of V1 (Richardson et al., 2005b); since the
V1 plasma instrument does not work, the MIR arrival times
at V1 were estimated by propagating V2 data to the distance of V1. These MIRs can affect the particle fluxes in
several different ways; as in the above example, they can
push the termination shock outward and sever or weaken
the spacecraft connection to the shock. Since the MIRs push
through the ambient solar wind, the magnetic field direction may change at the MIR or at the leading shock, thus
changing the magnetic connection between the termination
shock and the spacecraft, which changes the observed particle flux. Since MIRs are regions of enhanced magnetic field
strength and turbulence, energetic particles diffuse through
these regions slowly and pile up ahead of the enhanced
magnetic field region, forming intensity peaks as the particles are “snowplowed” ahead of the MIRs (McDonald et
al., 2000). Only one MIR has occurred since V2 entered the
sheath; it was associated with an increase of lower-energy
particles (28–540 keV) at the leading shock, probably due to
acceleration at the shock. The more-energetic 540–3500 keV
particles had intensity peaks after the shock; these could
result from snowplowing of particles ahead of the MIR. The
particle peak could also be behind the shock if the shock
has weakened so that higher-energy particles are no longer
accelerated; the particles already energized before the shock
source turned off would convect with the solar wind while
the shock would propagate ahead through the solar wind at
the fast mode speed (Decker et al., 2001).
Both V1 and V2 observe periodic enhancements in the
particle fluxes with periods of roughly a solar rotation. In
the first V1 foreshock encounter and in the V2 foreshock
453
data through mid-2007, these enhancements are not associated with changes in the plasma or magnetic field. However, when V1 was near the termination shock, in 2004,
many of these enhancements were associated with crossings of the heliospheric current sheet (HCS), but only from
negative to positive magnetic polarities (Richardson et al.,
2006b). One explanation for this correlation is that particles
from the termination shock current sheet drift (Burger and
Potgieter, 1989) along the HCS to V1, and at southern to
northern magnetic polarity HCS crossings the connection
to the termination shock is much closer. This mechanism
would give peaks once per solar rotation. This hypothesis
predicts that when V2, at southerly heliolatitudes, is close to
the termination shock, particle intensity peaks would be observed only at positive to negative polarity crossings of the
HCS. Particle peaks are observed by both V1 and V2 away
from the termination shock with periodicities of roughly a
solar rotation; the origin of these peaks is not understood.
6.
THE TERMINATION SHOCK CROSSING
The termination shock is where the solar wind plasma
makes its transition from supersonic to subsonic flow. The
solar wind plasma is compressed and heated and the magnetic
field strength increases. Zank (1999) reviews pre-encounter predictions of the nature and location of the termination
shock. Since the average spiral magnetic field upstream of
the termination shock is along the T axis (where R is outward from the Sun and T is in the solar equatorial plane and
positive in the direction of the Sun’s rotation), the termination shock is expected to be a quasi-perpendicular shock.
However, energetic particles accelerated or reaccelerated
at the termination shock and in the turbulent heliosheath
can propagate into the upstream region. When the Voyager
spacecraft are near the termination shock and connected to
it by the magnetic field, measured particle intensities and
anisotropies may be comparable to those at the termination
shock if pitch angle scattering is infrequent over the connection length and particles can traverse this length before
being convected shockward. Otherwise, intensities and anisotropies observed at Voyager will be reduced relative to
those at the termination shock.
The termination shock was crossed in December 2004 on
DOY 351, unfortunately during a tracking gap. The Voyager spacecraft are tracked by the DSN 10–12 hours/day
due to competition with other spacecraft. The day before the
shock crossing, the V1 Low Energy Charged Particle (LECP)
instrument saw an intense beam of field-aligned ions and
electrons at the same time as the Plasma Wave Spectrometer (PWS) experiment saw electron plasma oscillations;
theory predicts that electron beams will drive these waves
(Decker et al., 2005; Gurnett and Kurth, 2005). After the
shock crossing the lower-energy (tens of keV) ions jump in
intensity by about an order of magnitude above those in the
foreshock and the intensities become steady as shown in
Fig. 6. The termination shock particles shown in Fig. 7 are
well fit by power laws, consistent with diffusive shock ac-
454
The Solar System Beyond Neptune
Fig. 7. Typical spectra observed by Voyager in the heliosheath. The solid circles in the left panel are background-corrected ion (Z ≥ 1)
intensities from the LECP instrument and the open circles are from the Cosmic Ray Subsystem (CRS). The spectrum of the termination shock particles is a broken power law with break energy of ~3.5 MeV for protons and a helium. ACR helium dominates the midenergies in the right panel, with GCRs at higher energies in both panels. The predicted ACR spectra at the shock are shown for a
strong (r = 4) and a weak (r = 2.4) shock (Cummings et al., 2002).
celeration (Decker, 1988) theory for ions energized at a
shock. The spectra of termination shock particles both inside and outside the termination shock are best fit by two
power laws with a break energy at 3.5 MeV. The constancy
of the break energy suggests that the termination shock particles in the foreshock and heliosheath have the same source.
If the termination shock particles were accelerated by diffusive shock acceleration, then the shock strength can be
determined from the slope of the termination shock particle
power law and is 2.6 (2.4–3.0 with error bars) (Stone et al.,
2005). The shock strength is the compression ratio at the
shock (Ndownstream/Nupstream) and is 4 for a strong shock. An
alternative to diffusive shock acceleration has been suggested for the <20-MeV particles; a simple model propagating the measured pickup ions from Ulysses outward and
heating them at the termination shock can fit the observed
spectra at these energies (Gloeckler et al., 2005). In this
model, up to 80% of the solar wind dynamic pressure goes
into heating the pickup ions. A relatively weak shock is con-
sistent with the magnetic field observations, which showed
the magnetic field strength increased by about a factor of
3 at the shock and stayed high, convincing evidence that
V1 was in the heliosheath (Burlaga et al., 2005). The standard deviation of the magnetic field components increased
across the termination shock, consistent with predictions for
the heliosheath region, and the distribution of the field magnitudes changed from log normal in the solar wind to
Gaussian in the heliosheath. After V1 crossed the termination shock into the heliosheath, the intensities of ions of all
energies first increased, then fluctuated for 20–30 days, and
then became relatively steady, remaining essentially flat at
lower energies and increasing slowly at higher energies. At
the termination shock crossing, the intensity of 40–53-keV
ions increased tenfold, reached a peak, then dropped and
thereafter was relatively smooth, varying typically by no
more than ≈20% about a mean of ≈300 for the next six
months. The anisotropies went from beamlike in the foreshock to nearly isotropic in the heliosheath.
Richardson and Schwadron: The Limits of Our Solar System
The intensity of the more-energetic ACR particles increases at the shock but remains lower than the peak values
in the foreshock. As V1 moved deeper into the heliosheath
the ACR intensities increased. These ions are isotropic, not
flowing in field-aligned beams as in the foreshock. The major surprise at the shock crossing was that the ACR intensities did not peak and the ACR spectra did not unroll into
a power law distribution. For almost 30 years the paradigm
had been that ACRs were accelerated at the termination
shock through diffusive shock acceleration. This theory predicted a peak in ACRs at the shock and that the acceleration
process would result in a power law spectra at the termination shock. Instead, the ACRs intensities at the termination
shock did not reach even the level in the foreshock and the
ACR energy spectra barely changed. Figure 7 compares the
observed ACR spectra to those predicted for strong and
weak shocks. For energies less than 30 MeV the observed
intensities are far below those predicted. ACRs were not
being accelerated where V1 crossed the termination shock.
So where are they accelerated? That is still an open question. The intensities did increase further into the heliosheath.
Some investigators have suggested that the heliosheath itself
was the acceleration region.
Another attractive hypothesis is that the ACRs are accelerated at the flanks of the heliosphere where the connection time of a field line to the termination shock is longer
(McComas and Schwadron, 2006). Field lines first make
contact with a blunt termination shock near its nose. As
these field lines are dragged out into the heliosheath by the
solar wind, the connection point between the magnetic field
and termination shock moves toward the flanks. Anomalous
Fig. 8. The configuration of the termination shock and magnetic
field in the heliosheath from Schwadron and McComas (2007).
The bluntness of the termination shock leads to ACR acceleration on the distant flanks (McComas and Schwadron, 2006;
Schwadron and McComas, 2007). The red region shows the outer
modulation boundary of ACRs accelerated on the distant flanks
of the termination shock.
455
cosmic rays likely require substantial time to be accelerated; insufficient connection time near the nose could prevent acceleration to ACR energies. The much longer connection times toward the distant flanks of the shock may
provide the necessary acceleration time to achieve the very
high 100–300-MeV ACR energies (McComas and Schwadron, 2006; Schwadron and McComas, 2007). If the termination shock has a strong nose-to-tail asymmetry, then the
flanks would be connected by field lines that are well beyond V1 near the nose, as illustrated in Fig. 8.
Figure 9 documents the strong correlation between modulated ACR and GCR helium. The data points show V1 observations (McDonald et al., 2006) from January 2002 to February 2006; the upper panel, energy range 150–380 MeV/
nuc, shows GCRs, and the lower panels, energy ranges 10–
Fig. 9. The correlation between modulation of ACRs and GCRs.
The panels show the relative fluxes of ACRs and GCRs observed
by Voyager 1. The data (plus signs) were reduced by F. McDonald
and B. Heikkila. All fluxes have been normalized by their maximum values and the fluctuations over time (averaged over 26 days)
are plotted relative to one another. The respective energy ranges
and species are indicated on the axes. The solid curves show theoretical predictions based on a simple Parker-type diffusion-convection modulation model (see text for details).
456
The Solar System Beyond Neptune
21 MeV/n and 30–56 MeV/n, measure ACRs. The data are
normalized to the maximum value of the differential energy
flux in the respective energy ranges. Schwadron and McComas (2007) show how the strongly correlated modulation of ACRs and GCRs requires that the ACR accelerator
be on field lines that connect far beyond V1 on either the
distant reaches of termination shock flanks or far out in the
inner heliosheath near the heliopause.
Another possibility for explaining ACR acceleration in
the heliosheath is stochastic acceleration (Fisk, 2006). However, stochastic mechanisms encounter several problems
with explaining ACR acceleration at high energies: (1) the
rates of stochastic acceleration are typically lower than the
rate of diffusive acceleration at the termination shock, which
makes it difficult to generate highly energetic ACRs in the
one- to a few-year time frame before they become multiply
ionized (Jokipii, 1987; Mewaldt et al., 1996); and (2) the
mechanism must lead to ACRs generated only far out in the
heliosheath, or even toward the tail, whereas a distributed
process seems more likely to generate a very distributed
source. Although stochastic processes may not act quickly
enough to accelerate ACRs to high energies, they are critical for generating the suprathermal seed population, which
is naturally favored for diffusive acceleration at the termination shock (Fisk, 2006).
McDonald et al. (2006) suggest that the lack of an ACR
peak was a time-dependence effect. A MIR passed V1 and
encountered the termination shock just before the termination shock crossed V1; perhaps this reduced the acceleration efficiency and the ACR spectra recovered with time,
so the increase in intensity observed after the V1 termination shock crossing is a time, not radial distance, effect. This
scenario accounts for the larger ACR fluxes in the foreshock
than at the termination shock; the MIR reduces the ACR
fluxes at the termination shock from the previous, larger
intensities that were observed flowing into the foreshock
from the termination shock. Over six months later, the ACR
fluxes were still recovering from this large temporal change.
This hypothesis will be tested when V2 crosses the termination shock.
Another surprise was that after the termination shock
crossing V1 stayed in the same magnetic sector for 150 days
(Burlaga et al., 2005). Since the Sun rotates every 25.4 days
and the HCS is tilted, on average two HCS crossings should
occur every 25 days from southern magnetic polarities to
northern and back. These data were even more puzzling
since V1 was in the southern polarity sector, where V1
should spend less time since it is at northerly latitudes. Sector crossings did resume (Burlaga et al., 2006); the explanation for the initial lack of crossings is that the termination shock and heliosheath were moving inward at a speed
comparable to the outward heliosheath flow speed, so V1
remained in the same heliosheath plasma as it moved outward (Jokipii, 2005). The solar wind pressure was decreasing, which would make the termination shock and heliosheath move inward. The radial speeds deduced from the
LECP data range from –30 to 100 km/s but average 30–
50 km/s in this time period, consistent with the termination
shock moving inward.
An interesting feature in the magnetic field data are large
fluctuations in the field magnitude, a factor of 3, that are
purely compressional (the field direction does not change)
(Burlaga et al., 2006). These features look qualitatively like
the mirror mode waves observed in the magnetosheaths of
Jupiter and Saturn (Joy et al., 2006; Bavassano-Catteneo
et al., 1998); sometimes they are dips below the background
field value as are observed in low β regions of planetary
magnetosheaths, and sometimes they are increase above the
background field as observed in high β regions. The periods
of these waves are roughly an hour. If they are mirror mode
waves, they result from an instability that occurs in high β
plasmas when Tperp > Tpar, where Tperp and Tpar are the temperatures perpendicular and parallel to the magnetic field
direction. These conditions often occur behind perpendicular shocks such as the termination shock. As Voyager gets
further from the shock, one expects that the plasma will, as
a result of these waves, become more stable and the waves
less frequent.
7.
THE HELIOSHEATH
The spectra of the termination shock ions in the heliosheath are much less variable than in the foreshock. The
steadiness of the spectrum suggests that the spectral shape is
little affected by modulation in the heliosheath, as expected
(Jokipii and Giacalone, 2003). The energy spectrum in the
heliosheath hardens (has a larger percentage of high-energy
ions) with time; the intensities of the higher-energy ions
increase while those of the lower-energy ions are roughly
constant. The angular distributions of ions from 40 keV to
at least 30 MeV were highly anisotropic during the 2.5-year
approach of V1 to the termination shock. In the heliosheath
the ion anisotropies are much smaller and more variable in
direction (Decker et al., 2005; Krimigis et al., 2005).
Figure 6 shows that 5–6 days after the termination shock
crossing, V1 began measuring a positive radial flow component fluctuating around VR ≈ +100 km s–1, which continued for 27–28 days. Near 2005.045, VR decreased rapidly (within about 5 days) and flow became inward from
2005.063 to 2005.300. From about 2005.30 to 2005.4, VR
was 0 to +30 km s–1. VR increased again around 2005.45
and remained near +50 to +100 km s–1 until at least 2005.54.
These excursions of VR over a period ≈90–140 days are
consistent with V1 sampling the variable flow downstream
of an inwardly moving termination shock. Convective flow
at V1 that fluctuates about zero implies that V1 would essentially be riding with and sampling the same parcel of
plasma for an extended period. Tangential flows are also observed in the direction away from the heliospheric nose direction, as expected for flow around the heliosphere (Decker
et al., 2007).
The relationship between ACRs and termination shock
particles is unknown. Both are deficient in C ions, indicating
that they are accelerated pickup ions (Krimigis et al., 2003).
Richardson and Schwadron: The Limits of Our Solar System
However, the H/He ratio is ~10 for termination shock particles (Krimigis et al., 2003) and ~5 for ACRs (Cummings
et al., 2002), suggesting that He is more easily accelerated
to ACR energies than H (Zank et al., 2001).
The flow streamlines in the heliosheath were calculated
by Parker (1963), assuming subsonic, incompressible, irrotational flow. Using a similar flow model and a kinematic
approximation, Nerney et al. (1991) and Washimi and Tanaka (1996) calculated the variation of the magnetic field
B in the heliosheath. Close to the termination shock, the velocity remains nearly radial and B remains nearly azimuthal.
The speed decreases as the plasma moves toward the heliopause; consequently, the magnetic field strength B increases. Cranfill (1971) and Axford (1972) showed that if
the heliosheath is thick enough, B might become strong
enough to influence the flow. Nerney et al. (1993) estimated
that B might influence the flow in a restricted region. The
flow is deflected away from the radial direction as it approaches the nose of the heliopause, and B develops a radial
component. At the heliopause, B must be parallel to the surface (unless significant reconnection occurs at the boundary). The observations of B in the heliosheath show features not described by the simple kinematic models. While
B does have a strong azimuthal component, a persistent meridional component (northward) is observed that was not
predicted (Burlaga et al., 2005).
The magnetic field direction in the solar wind is often
alternately toward and away from the Sun for several days.
The existence of these “sectors” is related to extensions of
fields from the polar regions of the Sun to the latitude of the
observing spacecraft. Sectors have been observed by V1 out
to 94 AU (Burlaga et al., 2003b, 2005).
8.
THE HELIOPAUSE
As the Voyagers continue their trek through the uncharted heliosheath, the next milepost in their journey will
be the heliopause, which we think they will encounter between 2013 and 2021. The crossing of the heliopause will
be a truly historic occasion; the first manmade vehicle to
leave the solar system. The heliopause is thought to be a
tangential discontinuity with a large jump in the plasma
density and a rotation and change in the magnitude of the
magnetic field. Given the lack of an operational plasma instrument on V1, the first crossing of the heliopause will
probably be identified by a durable change in magnitude
and/or direction of the magnetic field. If the model asymmetries discussed above are real, then the V1 and V2 crossings could be close to each other in time.
The flow velocity vectors are expected to become tangent to the nominal heliopause surface. However, variations
of the heliopause surface and fluctuations in the flow velocities could make it difficult to use these data to identify
a heliopause crossing. In analogy with planetary magnetopauses, the heliopause is probably a complex surface that
varies locally in thickness and orientation and is the site of
patchy reconnection and a variety of plasma instabilities.
457
The appearance of large variations in flow velocities associated with disturbances from such relatively small-scale dynamical processes may indicate that V1 is near the heliopause. Reconnection at the heliopause-LISM interface could
accelerate low-energy ions (and electrons) and create large
departures from the normal heliosheath intensities and
anisotropies. Intermittent reconnection along a slightly corrugated heliopause surface might appear in the low-energy
ion (and possibly, electron) intensity-time profiles as gradual intensity increases with superposed, anisotropic intensity spikes as the spacecraft approached the heliopause, not
unlike structures encountered by V1 as it approached the
termination shock. These comments are speculative, drawing largely upon analogies to planetary magnetopauses;
however, as with the termination shock, energetic particle
observations may enable us to remotely sense the heliopause well before the Voyagers cross it.
Instabilities on the heliopause have been the subject of
theoretical conjecture (cf. Fahr, 1986; Florinski et al.,
2005). In principle, the Rayleigh-Taylor instability could be
important near the nose of the heliosphere where the difference between the flow densities are large. This instability
requires an effective “gravity” or destabilizing force. Liewer
et al. (1996) proposed that coupling between ions and neutrals via charge exchange would produce such a destabilizing force. Near the flanks, away from the nose, the KelvinHelmholtz instability is more likely to be important where
the shear velocity across the heliopause is large. Either type
of instability would produce motion of the heliopause superposed on the “breathing” expected to occur due to varying solar wind pressure. Estimates of the magnitude of the
resulting excursions are tens of AU and have periods on the
order of 100 years or more (Liewer et al., 1996; Wang and
Belcher, 1998). For Voyager, the long periods make it unlikely that such wave motion can be detected, but the additional radial motion from such instabilities, should they
exist, increases the uncertainty in the distance to the heliopause.
9.
THE HYDROGEN WALL
Plate 11 shows that a stagnation point forms in the flow
at the nose of the heliosphere. The plasma flow speed in
this region slows, so the plasma is compressed and heated.
Since the H+ and the H are coupled via charge exchange,
the H slows as well. Since the flux of H must be conserved,
slower H speeds translate to larger H densities (see Baranov
and Malama, 1993, 1995; Zank et al., 1996). Plate 11 shows
the region of denser H known as the hydrogen wall in front
of the heliosphere. This wall has been observed. Observations of H Lyman-α emission from nearby stars in the nose
direction show that the H absorption is red-shifted relative
to the other lines by 2.2 km/s; its width is too broad to be
consistent with the T = 5400 ± 500 K temperature suggested
by the width of the deuterium line (Linsky and Wood, 1996).
The profiles could only be fit assuming additional absorption by hot H, such as the compressed H in the hydrogen
458
The Solar System Beyond Neptune
wall. Two H components were needed to fit the data, one
from the heliospheric hydrogen wall (which, since the H is
moving sunward slower than the local interstellar cloud H,
produces absorption on the red side of the H line) and one
from the hydrogen wall of the nearby stars (which, since
the H is moving sunward faster than the local interstellar
cloud H, produces absorption on the blue side of the H line)
(Linsky and Wood, 1996; Gayley et al., 1997; Izmodenov
et al., 1999, 2002; Wood et al., 1996, 2000). Many hydrogen walls have been detected in front of other astrospheres,
allowing comparison to the heliosphere and providing validation to our general view of the heliosphere structure.
10.
THE HELIOSPHERIC RADIO EMISSIONS
The Voyager spacecraft periodically detect radio emission from the boundaries of our solar system. This radio
source is the most powerful in the solar system with an
emitted power of over 1013 W (Gurnett and Kurth, 1996).
Plate 12 shows a spectrogram of these emissions, plotting
intensity as a function of time and frequency. These waves
have two frequency components, a constant frequency 2kHz component and a component where the frequency rises
from 2.4 to 3.5 kHz over a period of 180 days. The two
largest events occurred in 1983–1984 and 1992–1994, a few
years after solar maxima. Large interplanetary shocks or
MIRs were hypothesized to trigger these emissions when
they entered the higher-density hydrogen wall beyond the
heliopause. But despite the occurrence of several large events
after the 2000 solar maximum, only very weak emission
was observed in 2002–2003, even though Voyager is presumably much closer to the radio source. The source region and trigger for these emissions is not well understood.
11.
THE HELIOSHEATH: THE LOCAL
INTERSTELLAR MEDIUM
The local interstellar medium surrounding the heliosphere is not a constant. As the Sun and interstellar clouds
move through the galaxy, the Sun has been and will be in
very different environments. The interstellar medium has
hot tenuous regions and cool dense regions that coexist in
pressure balance. The Sun is now in a hot tenuous region
of the interstellar medium known as the local bubble. This
region is about 200 pc across and has T = 106 K and an
average neutral H density of 0.07 cm–3. This density is about
10 times smaller than the average interstellar H density. The
local bubble was created by several supernova explosions
that occurred on the order of a million years ago, blowing
a hole in the interstellar medium. The local bubble contains
several clouds, including the local interstellar cloud in which
the Sun now resides, with an H density of about 0.2 cm–3.
The Sun is very near the edge of the local interstellar cloud
and will enter a different interstellar environment in less than
4000 years. It entered the current local interstellar cloud
within the past 40,000 years. Possible values of the densities
encountered by the heliosphere in the past and future are
0.005–15 cm–3 with Sun-cloud relative velocities of up to
100 km/s (Müller et al., 2006). Given possible values of the
interstellar pressure, in the past and future the heliopause
could be located anywhere between 12 and 400 AU and the
termination shock from between 11 and 250 AU. For a very
small heliosphere, cosmic-ray fluxes would dramatically increase, affecting the climate and increasing rates of genetic
mutation. When the heliosphere is in the local bubble, it
will be greatly expanded. The GCR fluxes at Earth will be
similar to present-day fluxes, but the fluxes of ACRS
and neutral H will be greatly reduced.
12.
FUTURE OBSERVATIONS
12.1. Voyager
The Voyager spacecraft continue their journeys, V1 moving out about 3.5 AU/year and V2 about 3 AU/year. Voyager 2 is in the termination shock foreshock region; models suggest that this region only extends 4–7 AU from the
termination shock. In June 2007, V2 had already been in
the foreshock region for almost 8 AU. The location of the
termination shock and foreshock change with the solar wind
pressure and thus the time spent in the foreshock can be
longer than its width would suggest, as when the termination shock was moving out just ahead of V1, but V2 will
probably cross the termination shock before the end of 2008.
This crossing will provide information on the asymmetry
of the heliosphere. It will also provide the first plasma data
from the heliosheath, providing direct information on the
plasma flow, shock compression, and heating at the termination shock. Figure 10 shows the V1 and V2 trajectories
in distance and heliolatitude. Model estimates of the locations of the termination shock and heliopause are shown
shown by the lines.
The V1 heliopause crossing is likely to occur between
2013 and 2021 and the V2 heliopause crossing should occur
in the same time frame. The spacecraft will have sufficient
Fig. 10. Trajectories of the Voyager spacecraft in distance and
heliolatitude. The range of predicted heliopause (for V1 and V2)
and termination shock (for V2) distances are shown by the solid
parts of the trajectory.
Richardson and Schwadron: The Limits of Our Solar System
power to operate all instruments until 2016; after this time,
power-sharing will extend the useful life of the spacecraft
beyond 2020. Thus the Voyagers are likely to provide the
first in situ measurements of the local interstellar cloud.
12.2.
New Horizons
New Horizons is a NASA mission that will fly by the
Pluto-Charon system and then proceed to other Kuiper belt
objects. Again, due purely to serendipity, Pluto is now located toward the nose of the heliosphere. Although this
mission will not return data from the interstellar medium,
it will measure the pickup ions directly, providing new information on the densities and composition of the interstellar neutrals entering the solar system.
12.3.
459
lar medium. For example, Gloeckler et al. (2005) extrapolated the V1 LECP proton spectra (40–4000 keV) upstream
and downstream of the termination shock to lower energies
and demonstrated that 80% of the upstream ram pressure
goes into heating the pickup protons. Thus V1 provides
the input proton spectrum for the global heliosheath (albeit
only at the V1 location) that will be imaged in IBEX hydrogen ENAs. And, returning the scientific favor, the very
first IBEX half-year images will, at a glance, reveal asymmetries in the large-scale heliosheath configuration that bear
on the puzzling directions of the magnetic field and energetic particle streaming observed by V1 and V2. The IBEX
all-sky images, in concert with local measurements from
Voyager, will provide a quantum leap forward in our understating of the outer heliosphere and its interstellar interaction.
Interstellar Boundary Explorer
12.4.
The Interstellar Boundary Explorer (IBEX) is an exciting new NASA mission that will observe the large-scale
structure of the heliosphere directly. It will be launched into
a highly eccentric Earth orbit in June 2008 [see the general description of the mission in McComas et al. (2004)].
All-sky energetic neutral atom (ENA) images of the global
structure of the termination shock and heliosheath will be
accumulated every 6 months. These ENAs are generated by
charge-exchange collisions between the inflowing interstellar neutral gas and the protons that are heated/accelerated
at the termination shock and that subsequently populate the
entire heliosheath. Two IBEX single-pixel telescopes will
image these hydrogen ENAs over semiannually precessing great circles in the sky at low energies (10 eV to 2 keV)
and in an overlapping band of higher energies (300 keV to
6 keV). The intensities of the ENAs depend strongly on both
the shape of the proton spectrum and the plasma convection flow pattern in the heliosheath. Where the anti-Sunward
component of the heliosheath convection velocity is high,
the ENA intensity directed toward Earth is low, and vice
versa. The ENA foreground from the heliosphere between
Earth and the termination shock is negligible because the
solar wind and pickup ion populations are being convected
away from the Sun at the solar wind velocity (which is
greater than the Earth-directed velocities of the thermal or
pickup ions). Beyond the termination shock, the shocked
decelerated flow in the heliosheath allows heated/accelerated protons to reach Earth and be imaged by IBEX.
The ENA emission seen from Earth is accumulated along
the entire line of sight outward through the inner heliosheath, across the heliopause, and into the outer heliosheath.
The measured ENA intensity integrates the effect of the termination shock heating/acceleration and heliosheath plasma
flow over several tens of AU, thus yielding diagnostics of
this immense boundary region (Gruntman et al., 2001).
Even though V1 and V2 do not directly measure protons
in the energy range corresponding to the IBEX ENA hydrogen, these complementary measurements will prove of
inestimable value in constraining the theoretical models of
the global interaction of the solar wind with the interstel-
The Interstellar Probe
Finally, in the era after Voyager and IBEX, we hope to
send a new, optimally instrumented mission called the Interstellar Probe out beyond the termination shock, heliopause, and bow shock and into the pristine interstellar medium. Such a mission will require propulsion capable of
accelerating a spacecraft to immense speeds if results are
to be gained in a reasonable mission lifetime. Recent studies
have examined both solar sails and nuclear-powered electric propulsion. While the results of such a mission are still
decades off, we need to begin soon if we are to finally send
a probe out into the material of the stars.
Acknowledgments. J.D.R. was supported under NASA contract 959203 from the Jet Propulsion Laboratory to the Massachusetts Institute of Technology. N.A.S. was supported by the NASA
Interstellar Boundary Explorer mission, which is part of the GSFC
Explorer Program, and by NASA’s Earth-Moon-Mars Radiation
Environment Model (EMMREM) Project.
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