the limits of our solar system

Richardson and Schwadron: The Limits of Our Solar System 443

443

The Limits of Our Solar System

John D. RichardsonMassachusetts Institute of Technology

Nathan A. SchwadronBoston University

The heliosphere is the bubble the Sun carves out of the interstellar medium. The solar windmoves outward at supersonic speeds and carries with it the Sun’s magnetic field. The interstellarmedium 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. Beforethese plasmas interact at the heliopause, they both go through shocks that cause the flows tobecome subsonic and change direction. Voyager 1, which recently crossed the termination shockof the solar wind, is the first spacecraft to observe the region of shocked solar wind called theheliosheath. The interstellar neutrals are not affected by the magnetic fields and penetrate deepinto the heliosphere, where they can be observed both directly and indirectly. This chapterdiscusses these outer boundaries of the solar system and the interactions between the solar windand the interstellar medium, and describes the current state of knowledge and future missionsthat may answer the many questions that remain.

1. INTRODUCTION

The limits of the solar system can be defined in manyways. The farthest object orbiting the Sun could be one. Thefarthest planet could be another. We define this limit as theboundary between the Sun’s solar wind and the interstellarmedium, the material between the stars in our galaxy. Theinterstellar medium is not uniform, but is composed of manyclouds with different densities, temperatures, magnetic fieldstrengths and directions, and flow speeds and directions.The Sun is now embedded in a region of the interstellar me-dium called the local interstellar cloud. The solar wind flowsoutward from the Sun and creates a bubble in the interstel-lar medium.

The solar wind pushes into the local interstellar clouduntil it too weak to push any farther, then turns and movesdownstream in the direction of the local interstellar cloudflow. The plasmas in these two winds cannot mix, since theyare confined to magnetic field lines and magnetic fields can-not move through one another. Neutrals, however, are notbound by magnetic fields and can cross between regions. Solocal interstellar cloud neutrals can and do penetrate deepinto the solar system. The boundary between the solar andinterstellar winds is called the heliopause.

Both the solar wind and the interstellar wind are super-sonic; thus a shock forms upstream of the heliopause in eachflow. At these shocks, the plasmas in each wind are com-pressed, heated, become subsonic, and change flow direc-tion so that the plasma can move around the heliopause (forinterstellar medium) or down the heliotail (for the solarwind). The shock in the solar wind is called the terminationshock. The shock in the local interstellar cloud is called thebow shock. The region between the termination shock and

the heliopause is the inner heliosheath and the region be-tween the bow shock and the heliopause is the outer helio-sheath.

Plate 11 shows the equatorial plane from a model simu-lation of the heliospheric system (Müller et al., 2006). Thecolor coding in the top panel gives the plasma temperatureand in the bottom panel shows the neutral H density. Theplasma flow lines are shown in the top panel. The trajecto-ries of Voyager 1 (V1) and Voyager 2 (V2) are also shown;V1 is headed nearer the heliospheric nose than V2. We notethat it is only by serendipity that the Voyager spacecraft areheaded toward the nose of the heliosphere; the Voyager tra-jectories were driven solely by the locations of the planetsthat were part of Voyager’s grand tour of the solar system.The main heliospheric boundaries, termination shock, helio-pause, and bow shock are labeled in the upper panel. Thelocal interstellar cloud is only barely supersonic, so the heat-ing that occurs at the bow shock is small; at this boundarythe flow of local interstellar cloud plasma begins to divertaround the heliosphere. In the solar wind the flow is radi-ally outward until the termination shock is reached. At thetermination shock, the plasma is strongly heated and theflow lines bend toward the tail. At the heliopause, the localinterstellar cloud and solar wind flow are parallel; plasmagenerally cannot cross this boundary. The neutral densityincreases at the nose of the bow shock, forming an enhanceddensity region known as the hydrogen wall, and decreasesat the heliopause.

We are just beginning to learn about these outer regionsof our heliosphere and how they interact with the local in-terstellar cloud. Until V1 crossed the termination shock in2004 at 94 astronomical units (AU) (1 AU is the distancefrom the Sun to Earth), we could only guess the boundary

444 The Solar System Beyond Neptune

locations. Now we have observed one termination shockcrossing and have several years of heliosheath data, but westill have much to learn about the other heliosphere bounda-ries and the global aspects of these boundaries. This chapterdiscusses what we know about the interaction of the helio-sphere with the local interstellar cloud and these boundaryregions, the unsolved problems, and the future observationswe hope will help answer these questions.

2. THE HELIOSPHERE

The heliosphere is the region dominated by plasma andmagnetic field from the Sun. No spacecraft has yet crossedoutside this region; estimates of the size of the heliosphereat the nose (in the upwind direction) range from 105 to150 AU (Opher et al., 2006). In the tail, the heliosphereextends many hundreds and probably thousands of AU. Thelocation of the nose of the heliopause is determined by pres-sure balance between the outward dynamic pressure of thesolar wind and the pressure of the local interstellar cloud,which has contributions from the dynamic pressure, thermalpressure, and magnetic pressure (see Belcher et al., 1993).The dynamic pressures are given by 1/2ρV2 where ρ is themass density and V is the speed. For the local interstellarcloud the magnetic pressure B2/8π and the thermal pressuresnkT could be important but the magnitudes are not wellconstrained. We discuss first the unperturbed flows from theSun and in the local interstellar cloud, then discuss the in-teractions of these two flows.

2.1. The Solar Wind

The Sun is a source of both photons and plasma; thefluxes of each fall off as 1/R2. The effects of solar radiationon the heliosphere/local interstellar cloud interaction aresmall compared to those of the solar wind. The first detec-tions of the solar wind in the early 1960s (Neugebauer andSnyder, 1962; Gringauz, 1961) confirmed Parker’s (1958)theoretical predictions of the basic solar wind structure. Thesolar wind has been thoroughly measured and monitoredby spacecraft since that time. In addition to the near-equa-torial spacecraft in the inner heliosphere, Pioneers 10 and11 and V1 and V2 have observed the solar wind in the outerheliosphere, in the case of V1 past 100 AU. Ulysses has pro-vided measurements in the high-latitude heliosphere. Theseobservations give a good picture of the solar wind parame-ters throughout the heliosphere.

The solar wind is a supersonic flow of mostly protons,with on average 3–4% He++ (this percentage varies over asolar cycle) (Aellig et al., 2001) and a small amount ofhighly charged heavy ions. The Sun has a magnetic fieldthat reverses polarity every 11-year solar cycle. To first or-der, plasma cannot cross magnetic field lines, so the solarwind outflow drags the solar magnetic field lines outward.This outward motion combined with the Sun’s 27-day rota-tion causes the magnetic field lines to form spirals; at Earththe spiral angle is about 45°. In the outer heliosphere these

spirals become tightly wound and the magnetic field angleis nearly 90° (the field is almost purely tangential). Sincethe magnetic field lines from the northern and the southernhemispheres of the Sun have opposite polarity, a helio-spheric current sheet (HCS) forms near the equator, acrosswhich the magnetic field direction reverses. The HCS passesthrough the termination shock and has been observed in theheliosheath past 100 AU. The average solar wind speed atEarth is about 440 km/s with a density of about 7.5 cm–3,an average temperature of 96,000 K, and an average mag-netic field strength of about 5 nT; these values are basedon a solar cycle of data from the Wind spacecraft from 1994to 2005. The average dynamic pressure is about 2.25 nP;this pressure determines how big a cavity the solar windcarves into the local interstellar cloud.

2.2. The Local Interstellar Cloud

The local interstellar cloud surrounds the heliosphere. Itcontains neutral and ionized low-energy particles, H, He, O,H+, He+, and O+, as well as very-high-energy galactic cosmicrays (GCRs). These GCRs are observed at Earth but theirintensities are modulated (reduced) by the solar wind mag-netic field as they move through the heliosphere. Plate 11shows how the low-energy plasma deflects around the helio-sphere. Helium is the only major species that makes it toEarth in pristine form and thus provides the best determi-nation of the properties of the interstellar medium. Obser-vations of the H and He emission combined with modelsof the neutral flow give estimates of the neutral speed, den-sity, and temperature as well as the flow direction.

Hydrogen is the largest component of the interstellar me-dium, but it is difficult to determine the local interstellarcloud properties from H observations because a large frac-tion of the H charge exchanges with the local interstellarcloud H+ before it enters the heliosphere. Charge exchangeoccurs when an ion and a neutral collide and an electrongoes from the neutral to the ion. The former ion is a nowneutral and is not bound by magnetic fields, so it continuesto move at the speed it had when it was an ion. The formerneutral, now an ion, is bound to the magnetic field andaccelerated to the speed of the bulk plasma flow. Outsidethe bow shock the plasma and neutrals have the same speedand temperature so charge exchange has no effect. But in-side the bow shock, the plasma parameters are differentfrom those in the interstellar medium since the plasma isheated, compressed, and deflected at the shock. The plasmaslows further as it piles up in front of the heliosphere. Sowhen charge exchange occurs, the new neutrals have speedsand temperatures from the shocked plasma, not the pristineinterstellar medium. The H observed in the heliosphere hasan inflow speed of only 20 km/s, 6 km/s less than the inter-stellar medium speed (Lallement et al., 1993), confirmingthat the H has charge exchanged with local interstellar cloudH+, which has slowed down near the heliosphere boundary.The plasma flow in this region has begun to divert aroundthe heliosphere, so the flow direction of the H is also altered.


Observations with the SOHO/SWAN instrument show thatthe arrival direction of the interstellar H differs from thatof the interstellar He by about 4° (Lallement et al., 2005).This offset results from charge exchange of the H with thedeflected H+ in front of the heliosphere. The direction ofthe offset tells us the direction of the plasma flow in thelocal interstellar cloud just outside the heliopause and thusprovides information on the magnetic field direction (Lalle-ment et al., 2005). The magnetic field strength in the localinterstellar cloud is not known; generally it is thought tobe on the order of 1 µG with an upper limit of 3–4 µG.

The interstellar neutrals enter the heliosphere and absorband emit sunlight at specific wavelengths. Backscattered HLyman-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 theSun, they are accelerated inward by gravity but pushed out-ward 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 andis accelerated inward. Hydrogen also has a much shorterionization time in the solar wind than He. The result is thatmost H becomes ionized before it reaches 5–10 AU. Thepeak in backscattered Lyman-α intensity is roughly in thedirection of the local interstellar cloud flow with a smalldeflection occurring from the charge exchange in the hy-drogen wall.

The ionization time of He is long; most He is not ionizeduntil it is inside 1 AU. The He with trajectories that remainoutside 1 AU survive and flow downstream out of the helio-sphere. The Sun’s gravity bends the trajectories so that theHe is focused downstream of the Sun, where the maximumemission from He is observed (Axford, 1972; Fahr and Lay,1973). The location of the maximum He emission is thusopposite to the direction of the incoming He flow. Obser-vation of the direction with peak He intensity is one methodused to determine the local interstellar cloud flow direction.

The interstellar neutrals are ionized in the solar wind viacharge exchange with protons (the dominate process for H)or photoionization. These neutrals can be observed directlyas H+ and He+ pickup ions. He+ is rarely observed in thesolar wind and thus is a good marker of interstellar neutrals.Both H+ and He+ pickup ions have distinctive signatures inthe solar wind; when the neutral H and He become ionizedthey 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 thesolar wind speed, so these ions show a sharp energy cutoffat four times the solar wind energy.

The energy to accelerate the pickup ions comes from thebulk flow of the solar wind, so an indirect way to measurethe pickup ion density (discussed further below) is to look athow much the solar wind slows down as it moves outward.These direct and indirect observations can be combined withmodels of the neutral inflow and loss to estimate speeds,temperatures, and densities of H and He in the interstellarmedium.

The gas detector on Ulysses measures the neutral Hevelocity and temperature directly (Witte, 2004; Witte et al.,1992). Even with these direct measurements, determinationof the density of the local interstellar cloud still requiresmodeling since some of the He is lost on its way through theheliosphere. All these measurements taken together providefairly tight constraints on the local interstellar cloud He neu-tral speed and temperature (Möbius et al., 2004).

The He moves toward the Sun from an ecliptic longitudeof 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 densityof He is 0.015 ± 0.002 cm–3. The collision timescales in thelocal interstellar cloud are less than the neutral and plasmalifetimes, so the plasma and neutral species should all havethe same speed and temperature far from the heliosphere.Since the H is altered by charge exchange with the plasmanear the heliosphere, the H density is more difficult to de-termine. Best estimates are on the order of 0.20 ± 0.03 cm–3

in the local interstellar cloud. Since the ionization fractiondepends on the electron temperature, which is not wellknown, the plasma density is also uncertain, with best esti-mates of 0.06–0.10 cm–3 based on the speed of the H meas-ured at Earth (Lallement et al., 1993).

3. THE TERMINATION SHOCK LOCATION

Although some properties of the local interstellar cloudare well constrained, others are not, which makes determin-ing the heliosphere size and shape difficult. We have ob-served one boundary location; the termination shock at theposition of V1 was at 94 AU on December 16, 2004 (Stoneet al., 2005; Decker et al., 2005; Burlaga et al., 2005). Thiscrossing provides one point with which to calibrate ourmodels. Since the average solar wind dynamic pressure is2.25 nP at 1 AU, one can roughly estimate that the totallocal interstellar cloud pressure must be 2.5 × 10–4 nP. Butthere are problems with this simple approach, which we de-scribe below.

3.1. Solar Wind Time Dependence

The average solar wind values given above are onlya piece of the story. The solar wind pressure is variableon scales from hours to 11-year solar cycles and likely haslonger-term variations. Small-scale variations of the pres-sure do not affect the boundary locations, but larger-scalevariations change the balance between the solar wind andlocal interstellar cloud pressures and cause the terminationshock and heliopause to move in and out. The solar winddynamic pressure changes by a factor of about 2 over a solarcycle, with minimum pressure near solar maximum (whensunspot numbers are highest), then a rise in pressure witha peak a few years after solar maximum, followed by a de-crease to the next solar maximum (Lazarus and McNutt,1990). Models of the heliosphere show that these pressurechanges drive an inward and outward “breathing” of the he-liosphere; the termination shock location changes by about


10 AU over a solar cycle and the heliopause by 3–4 AU(Karmesin et al., 1995; Wang and Belcher, 1998). The solarwind structure also changes over the solar cycle. At solarminimum, the solar wind near the equator is slow (400 km/s) and dense (8 cm–3). The solar wind at higher latitudescomes from polar coronal holes and is fast, 700–800 km/swith a density of 3–4 cm–3. At solar maximum, the polarcoronal holes are not present and most of the flow is low-speed solar wind. Despite these latitudinal changes in thesolar wind, the dynamic pressure changes over the solar cy-cle are the same at all heliolatitudes (Richardson and Wang,1999), so the heliosphere is not expected to change shapeover a solar cycle.

Large events on the Sun can also affect the terminationshock position. Active regions on the Sun generate coronalmass ejections (CMEs), explosive discharges of materialinto the solar wind where they are called interplanetaryCMEs (ICMEs). ICMEs propagate through the heliosphere.When a series of ICMEs occurs, the later ones catch upto and merge with previous ICMEs, causing a buildup ofplasma and magnetic field known as a merged interactionregion (MIR) (Burlaga et al., 1984; Burlaga, 1995). Theseregions are usually preceded by a fast forward shock atwhich the speed, density, temperature, and magnetic fieldstrength all increase. Shocks are rare in the outer helio-sphere; most can be related to a large CME or a series ofCMEs on the Sun (see Richardson et al., 2005a). The dy-namic pressure in the MIRs can be a factor of 10 higherthan that in the ambient solar wind and these structures last1–3 months (Richardson et al., 2003). The MIRs drive thetermination shock outward several AU for 2–3 months, afterwhich the termination shock rebounds inward (Zank andMüller, 2003).

3.2. Local Interstellar Cloud Time Dependence

Changes in the pressure or magnetic field of the localinterstellar cloud could also cause the boundary locations tomove. The scale length for changes in the local interstellarcloud is not known. The Sun is very near the local inter-stellar cloud boundary and will cross into another interstel-lar medium environment in a few thousand years, and scalelengths for variations may be smaller near this boundary(Müller et al., 2006). Observations of He by Ulysses areconsistent with no change in the interstellar neutral densitysince the launch of Ulysses in 1990 (Witte, 2004).

3.3. Local Interstellar Cloud Effectson the Solar Wind

If we assume that the local interstellar cloud pressurewere constant, we can use the observed solar wind pres-sures to model the heliosphere boundary changes. Modelsmust include the effects of the local interstellar cloud neu-trals on the solar wind. These neutrals are the largest (bymass) component of the interstellar space outside about

10 AU. Figure 1 shows the density of plasma and neutralsvs. distance (we know now the termination shock and he-liopause in this figure are located too close to the Sun). Thesolar wind plasma falls off as R–2, then increases abruptlyat the termination shock and heliopause. The H density isabout 0.1 at the termination shock and 0.2 in the local in-terstellar cloud, then decreases rapidly inside 10 AU, theionization cavity inside which most H is ionized. Both theH and H+ increase between the bow shock and heliopauseand form the hydrogen wall. The pickup ion density/thermalion density ratio increases to the termination shock, whereabout 20% of the ions are pickup ions. This ratio is consis-tent with the observed slowdown of the solar wind. Outsideabout 30 AU the pickup ions (of local interstellar cloud ori-gin) dominate the plasma thermal pressure. So even thoughthe Sun has carved out its own space in the local interstel-lar cloud, the local interstellar cloud influence is very strong.

3.3.1. Solar wind slowdown. The neutral interstellar His continuously being lost, mostly by charge exchange withsolar wind protons. The neutrals formed from the solar windprotons escape the solar system at the solar wind speed. Thenew protons formed from ionized local interstellar cloudneutrals are accelerated from their initial speed of 20 km/sinward to the solar wind speed, 400–700 km/s outward. Asmentioned earlier, plasma does not flow across field lines;ions at speeds different from the bulk solar wind speed areaccelerated by an electric field E = –VXB. The initial ther-mal energy of the pickup ions is equal to the solar wind en-

Fig. 1. The density of the plasma and neutral components from1 to 1000 AU. The solar wind ions come from the Sun. The pickupions are interstellar neutrals that have been ionized in the solarwind. The densities of the solar wind and of the pickup ions jumpat the termination shock. Outside the heliopause, the ions are partof the local interstellar cloud. Both the ion and neutral densityincrease 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.


ergy, about 1 keV. The energy for the acceleration and heat-ing comes from the bulk flow of the solar wind.

Two consequences of these pickup ions are that the so-lar wind is heated and slows down. The amount of the slow-down depends on the density of the local interstellar cloudH, so measurements of the slowdown give an estimate of thelocal interstellar cloud density. This slowdown can be dif-ficult to measure because the solar wind speed varies withtime and solar latitude, as well as with distance. In the in-ner heliosphere, the slowdown is small and difficult to de-tect. In the outer heliosphere, determining the slowdownfrom the inner heliosphere to V2 requires (1) another space-craft in the inner heliosphere at a similar heliolatitude toprovide a speed baseline and (2) a model to propagate thesolar wind outward to determine how much of the speedchange comes from radial evolution of the solar wind andhow much is due to pickup ions. Figure 2 shows a compi-lation of various determinations of the slowdown. The plotshows the solar wind speed normalized to 1 AU. The hori-zontal lines show the average amount of slowdown; thelength of these lines shows the radial range covered in eachstudy. The vertical spread shows the uncertainties in theslowdown determination. The curves shows the predictedsolar wind slowdown for densities of interstellar H at thetermination shock of 0.07, 0.09, and 0.11 cm–3. The dataare 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 thatthe H density at the termination shock is not the density inthe local interstellar cloud; much filtration, or H loss, oc-curs between the bow shock and the termination shock.

By 95 AU, the solar wind speed has decreased by about20% [and since the mass flux, the speed (v) times the massdensity (ρ), is a constant, the density must increase by 20%].So the dynamic pressure, ρv2, is only 77% of what it was at1 AU. The effect of the interstellar H is to lower the solarwind dynamic pressure. This reduced pressure causes theheliosphere boundaries to be closer to the Sun. The solarwind slowdown is related to the pickup ion density by npu/nsw = (3γ – 1)/(2γ – 1)(δv/v0), where nsw is the solar winddensity, npu is the pickup ion density, δv is the solar windslowdown, v0 is the solar wind speed at 1 AU, and is theratio 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 movesout, it expands and thus might be expected to cool. The solarwind temperature does decrease out to about 30 AU, butnot as quickly as adiabatic cooling would predict, then in-creases with distance with some solar cycle variations su-perposed. Internally driven processes, such as shocks anddissipation of magnetic fluctuations, provide some heating,but not enough to produce the observed temperature pro-file (Richardson and Smith, 2003; Smith et al., 2001, 2006).The source of this heating is the pickup ions. These ionsdominate the thermal energy of the plasma outside about30 AU, and a small fraction of their energy is transferredto the thermal plasma. The mechanism for this energy trans-fer is through waves; the pickup ions initially have ringdistributions, which are unstable. Low-frequency waves aregenerated that isotropize the pickup ions, but about 4% ofthe 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 ofenergy transfered to the thermal ions are proportional to thesolar wind speed, the solar wind speed and temperature arestrongly correlated.

3.4. Termination Shock Model Results

Many models developed to determine the terminationshock motion, the asymmetry of the boundaries, and the dis-tance 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 thatincludes the effect of the local interstellar cloud neutrals onthe solar wind and use it to track the termination shock loca-tion. Figure 3 shows predictions of the termination shocklocation based on a two-dimensional model that includes theeffect of the local interstellar cloud neutrals on the solarwind. The model input is the solar wind pressure observedby V2; the profile is normalized to the V1 termination shockcrossing (Richardson et al., 2006a). The model shows thatV1 and the shock were both moving radially outward start-

Fig. 2. The plot shows the slowdown of the solar wind causedby the interstellar neutrals. The fraction of the initial solar windspeed is plotted vs. distance. The horizontal lines with error barsshow observed speed decreases; the lengths of these lines showthe radial width over which the slowdown was observed. Thevertical lines are error bars. The curves shows model predictionsof the solar wind speed as a function of distance for interstellarneutral H densities of 0.09 cm–3 (solid line), 0.07 cm–3 (upperdashed line), and 0.11 cm–3 (lower dashed line).


ing in 1999, that V1 made some close encounters to theshock, but that the shock crossing did not occur until thetermination shock moved inward in 2004.

3.5. Heliospheric Asymmetries

At Earth, the 45° average angle of the solar magneticfield to the solar wind flow direction results in asymmetriesin the magnetosheath plasma (Paularena et al., 2001) andthe magnetopause locations (Dmitriev et al., 2004). Thedraping of magnetic field lines around the magnetosphereresults in different field strengths on the dawn and dusksides, causing these asymmetries. The interstellar magneticfield could result in a similar distortions of the terminationshock and heliopause and significantly affect the width ofthe heliosheath (see, e.g., Linde et al., 1998; Ratkiewicz etal., 1998; Pogorelov et al., 2004). Voyager 2 entered the ter-mination shock foreshock region much closer to the Sun(75 AU) than V1 (85 AU) (Stone et al., 2005). This large adifference probably does not arise entirely from solar winddynamic pressure changes but also from asymmetries in theboundary locations, with the boundaries closer in the direc-tion of V2.

The nature of the heliosphere distortion depends on theorientation and strength of the interstellar magnetic field.Several suggestions for the magnetic field direction in thelocal interstellar cloud near the heliosphere have been pub-lished. The location of the sources of the heliospheric radioemissions (described below) lies along a line roughly par-allel to the galactic plane. Gurnett et al. (2006) suggest thatthe 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 thegalactic plane. A field orientation 60° from the galacticplane is derived from differences in the flow velocities ofinterstellar H and He (Lallement et al., 2005; Izmodenov etal., 2005).

Magnetohydrodynamic (MHD) models of the helio-sphere that include the effects of an interstellar magneticfield reveal that the distortion of the heliosphere is sensitiveto the angle between the field and the velocity of the inter-stellar wind. Opher et al. (2007) show that a field tilted 60°–90° from the galactic plane is needed to match the flow dataand the radio emission data. This tilt creates a large asym-metry in the model heliosphere in both the north/south andeast/west directions. The models show that magnetic fieldlines that intersect the termination shock extend 2 AU in-side the termination shock at the position of V1 and 4–7 AUinside the termination shock in the direction of V2 (Opher etal., 2006). These distances represent the average thicknessof the foreshock region in front of the termination shock.Small-scale magnetic field variations will affect these dis-tances locally. This model also predicts that the distance ofthe termination shock in the direction of V2 is 8–11 AUcloser than in the direction of V1 and the heliopause dis-tance is 13–24 AU closer in the direction of V2. This modeldid not include neutrals, which may reduce the asymmetry(Pogorolev et al., 2006). Other models also give asymme-tries with ratios of the termination shock distances at V1and V2 varying from 1.05 to 1.33. The predicted heliopauselocations vary from 145–157 AU at V1 and 109–139 AU atV2, with heliosheath thicknesses of 55–67 AU toward V1and 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 par-ticular, the effects of the HCS may be important, since thelocation of the HCS at the termination shock varies over asolar rotation, but this small-scale structure is hard to in-clude in a global model. Although the models differ quan-titatively, the results suggest that the average shock distanceat V2 could be ~5–10 AU closer than at V1. In addition,the solar cycle variation of the solar wind dynamic pressureis such that in 2007 the shock should be several AU closerthan average, suggesting that V2 might cross the shock in2007–2008. The models also indicate that the heliosheathis likely narrower at V2, so V2 may cross the heliopause atabout the same time as V1. Voyager 1 and V2 observationsof the asymmetry of the termination shock and heliosheathshould make possible improved models and estimates of thedistances to the heliopause.

4. COSMIC RAYS

Cosmic rays are high-energy charged particles that bom-bard Earth from above the atmosphere. These particles haveimportant effects at Earth. Cloud cover is directly correlatedwith cosmic ray flux, so these particles may affect Earth’sweather (Svensmark et al., 2007). Several thousand cosmic

Fig. 3. The predicted position of the termination shock using V2observations of the solar wind as input to a two-dimensional time-dependent hydrodynamic model. The termination shock locationsare extended into the future by repeating the pressure profile fromthe previous solar cycle. The V1 and V2 trajectories are shownby the dashed diagonal lines. The lower curve shows the termi-nation shock position shifted inward 10 AU to simulate a possibletermination shock asymmetry. The termination shock moved in-ward from 1995 to 1999, then moved outward ahead of V1 until2004, when it moved inward and crossed V1. The model suggestsV2 could cross the termination shock before the end of 2007.


rays pass through a person’s body every minute. These cancause biologic damage but also cause mutations that accel-erate evolution. A new source of cosmic rays discussed be-low is generated from the grains in the Kuiper belt (Schwad-ron, 2002). These cosmic rays were likely more intense inthe past and the termination shock was at times much closerto the Sun than it is today. Both of these effects may havecaused periods when Earth was ensconced in much moresevere radiation environments than at present. Several typesof cosmic rays come from the outer heliosphere, as dis-cussed below.

4.1. Galactic Cosmic Rays

Galactic cosmic rays (GCRs) possess energies above100 MeV and are thought to be accelerated in the shockwaves associated with supernova in our galaxy. The GCRsare predominately fully stripped nuclei but lesser fluxesof GCR electrons are also observed. Galactic cosmic raysrepresenting most elements in the periodic table have beendetected. Since cosmic rays are charged, they are affectedby magnetic fields in interplanetary space, the heliosphere,and near Earth. Although GCRs do reach Earth, many aredeflected by the heliospheric magnetic field before theyreach the inner solar system. Fewer GCRs are observed atEarth at solar maximum, when the solar wind magnetic fieldis stronger and more turbulent. Large merged interactionregions (MIRs) temporarily reduce the flux of GCRs asthese structures contain large and turbulent magnetic fields.An empirical relation between GCRs and the magnitude ofB shows that the GCR flux increases when B is above nor-mal and increases when B is below normal; this relationshipholds out to the termination shock (Burlaga et al., 2003a,2005). Spacecraft in the outer solar system see increasedGCR fluxes. These particles move so fast and are scatteredso thoroughly by the interstellar magnetic field that theirflux is likely to be fairly uniform outside the heliosphere.

4.2. Anomalous Cosmic Rays

Early cosmic-ray observers discovered an unusual subsetof cosmic rays that consisted of singly ionized ions (insteadof fully stripped nuclei) with energies of 1–50 MeV/nuc (seeMewaldt 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 werethought to arise only from neutral atoms in the interstellarmedium (Fisk et al., 1974) that drift freely into the helio-sphere through a process that has four essential steps: first,the neutral particles from the local interstellar cloud enter theheliosphere; second, the neutrals are converted into pickupions; third, the pickup ions are preaccelerated by shocks andwaves in the solar wind [see also section 5 and Schwadronet al. (1998)]; and finally, they are accelerated to their finalenergies at or beyond the termination shock (Pesses et al.,1981). Easily ionized elements such as C, Si, and Fe areexpected to be strongly depleted in ACRs, since these ele-

ments are not neutral in the interstellar medium and there-fore cannot drift into the heliosphere.

Instruments such as the Solar Wind Ion CompositionSpectrometer (SWICS) on Ulysses are able to detect pickupions directly. Ongoing measurements by cosmic-ray detec-tors on Voyager and Wind have detected additional ACRcomponents (Reames, 1999; Mazur et al., 2000; Cummingset al., 2002). We now understand that, in addition to thetraditional interstellar source, grains produce pickup ionsthroughout the heliosphere: Grains near the Sun producean “inner source” of pickup ions, and grains from the Kui-per belt provide an “outer” source of pickup ions and anom-alous cosmic rays.

Recent observations call into question not only thesources of ACRs, but also the means by which they are ac-celerated. The prevailing theory, until V1 crossed the ter-mination shock, was that pickup ions were energized at thetermination shock to the 10–100-MeV energies observed(Pesses et al., 1981). When V1 crossed the terminationshock, it did not see a peak in the ACR intensity as this the-ory predicts (Stone et al., 2005). Instead, the ACR intensi-ties continued to increase in the heliosheath. One suggestionis that the ACRs are accelerated in the flanks of the helio-sheath and the source region is not observed by V1 near thenose (McComas and Schwadron, 2006). Another sugges-tion is that the acceleration region was affected by a seriesof MIRs before the V1 crossing that diminished the accel-eration process (McDonald et al., 2006; Florinski and Zank,2006), in which case V2 may see a different ACR profilewhen it crosses the termination shock.

4.3. The Interstellar and Inner Source Pickup Ions

The telltale signature of interstellar pickup ions, as seenfor H+ in Fig. 4, is a cutoff in the distribution at ion speedstwice that of the solar wind in the spacecraft referenceframe. This cutoff is produced since interstellar ions are ini-tially nearly stationary in the spacecraft reference frame andsubsequently change direction, but not energy, in this framedue to wave-particle interactions and gyration, causing themto be distributed over a range of speeds between zero andtwice 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 dis-tributions do not cut off at twice the solar wind speed andhave a peak, contrary to the other ACR ions, which haveflat distribution. These ions are singly charged, which pre-cludes them from having been emitted directly by the Sun;if this were the case, they would be much more highlycharged, e.g., C5+ and O6+ are each common solar windheavy ions. The C+ and O+ distributions in Fig. 4 are consis-tent with a pickup ion source close to the Sun. As the ionstravel out in the solar wind, they cool adiabatically in thesolar wind reference frame, and instead of having a distri-bution that cuts off at twice the solar wind speed in thespacecraft reference frame, they have a peak near the solarwind speed. The solid curve that goes through the C+ and


O+ data points is based on a transport model (Schwadron,1998; Schwadron et al., 2000) for ions picked up close tothe Sun, thereby proving the source is in the inner solarsystem. Data like those presented in Fig. 4 suggested thatanother pickup ion source was present in addition to inter-stellar neutrals (Geiss et al., 1995; Gloeckler and Geiss,1998; 2000a; Schwadron et al., 2000). Since the source ispeaked near the Sun, as is the distribution of interplanetarygrains, it was hypothesized that grains were associated withthe source.

The initial, naive expectation was that the inner sourcecomposition would resemble that of the grains, i.e., en-hanced in C and O with strong depletions of noble elementssuch as Ne. However, the composition strongly resembledthat of the solar wind (Gloeckler et al., 2000a). This conun-drum was resolved by assuming a production mechanismwhereby solar wind ions become embedded within grainsand subsequently reemitted as neutrals. Remarkably, theexistence of the inner source was hypothesized many yearsprior to its discovery (Banks, 1971).

The concept that the inner source is generated due toneutralization of solar wind requires that sputtered atomsdo not strongly contribute to the source. However, for grainslarger than a micrometer, sputtering yields of the grainsare much larger than the yields of neutralized solar wind(Wimmer-Schweingruber and Bochsler, 2003). This suggeststhat the grains that give rise to the inner source are ex-tremely small (hundreds of angstroms). A small grain popu-lation generated through catastrophic collisions of largerinterplanetary grains would also yield a very large fillingfactor (the filling factor is the net area of the sky filled in by

grains divided by the net area of the sky over which they aredistributed; so, for example, a large solid object would havea filling factor of 1, whereas a finite set of points would havea filling factor of 0). A large filling factor is consistent withobservations of the inner source that require a net geomet-ric cross-section typically 100 times larger than that inferredfrom zodiacal light observations (Schwadron et al., 2000).The very small grains act effectively as ultrathin foils forneutralizing solar wind (Funsten et al., 1993), but are noteffective for scattering light owing to their very small size.The implications of this suggestion are striking: The innersource may come from a large population of very smallgrains near the Sun generated through catastrophic colli-sions of larger interplanetary grains (Wimmer-Schweingruberand Bochsler, 2003).

4.4. Pickup Ions from Cometary Tails

Comets are also a source of pickup ions in the helio-sphere. Until recently, it was thought that the only way toobserve cometary pickup ions was to send spacecraft tosample cometary matter directly. However, Ulysses has nowhad 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). Thesethree unplanned crossings of cometary tails by Ulysses sug-gest an entirely new way to observe comets through detec-tion of distant cometary tails.

4.5. Outer Source Pickup Ions andAnomalous Cosmic Rays

Recent observations from the Voyager and Wind space-craft resolved ACR components comprised of easily ion-ized elements [such as Si, C, Mg, S, and Fe (Reames, 1999;Mazur et al., 2000; Cummings et al., 2002)]. An interstellarsource for these “additional” ACRs, other than a possibleinterstellar contribution to C, is not possible (Cummings etal., 2002). Thus, the source for these ACRs must residewithin the heliosphere. A number of potential ACR sourcesare present within the heliosphere (Schwadron et al., 2002).The solar wind particles are easily ruled out as a sourcesince they are highly ionized, whereas ACRs are predomi-nantly singly charged. Discrete sources such as planets arealso easily ruled out since their source rate is not sufficientto generate the needed amounts of pickup ions. Another po-tential source is comets, but the net cometary source rate issufficiently large only inside 1.5 AU, a location so close tothe Sun that the generated pickup ions would be stronglycooled by the time they reach the termination shock, makingacceleration to ACR energies extremely difficult. Moreover,a cometary source would naturally be rich in C, which isnot consistent with the compositional observations of easilyionized ACRs.

The inner source (discussed in the preceding section) isanother possibility, but again, adiabatic cooling poses a sig-

Fig. 4. Observed distributions from Ulysses/SWICS of C+ (solidtriangles), O+ (open circles), and H+ (open triangles) vs. ion speedin the spacecraft frame normalized by the solar wind speed(Schwadron et al., 2000). The observations are compared withsimulated distributions of solar wind H+ (dashed black curve),interstellar H+ (upper gray dash-dot curve), inner source H+ (up-per thick black line), inner source C+, O+ (lower thick black line),and interstellar O+ (lower gray dash-dot curve).


nificant problem since these particles are picked up so closeto the Sun. It has been suggested that the inner source maybe substantially preaccelerated in the inner heliosphere,thereby overcoming the effects of adiabatic cooling. Con-trary to this suggestion, inner source observations in slowsolar wind show pronounced effects of adiabatic cooling(Schwadron et al., 1999) and charge-state observations ofenergetic particles near 1 AU indicate no evidence for accel-eration of the inner source (Mazur et al., 2002). The addi-tional population of ACRs requires a large source of pickupions inside the heliosphere and outside 1 AU.

Schwadron et al. (2002) suggest that the source is ma-terial extracted from the Kuiper belt through a series ofprocesses shown schematically by the lines in Fig. 5: First,micrometer-sized grains are produced due to collisions ofobjects within the Kuiper belt; grains spiral in toward theSun due to the Poynting-Robertson effect; neutral atoms areproduced by sputtering and are converted into pickup ionswhen they become ionized; the pickup ions are transportedby the solar wind to the termination shock and, as they areconvected, are preaccelerated due to interaction with shocks

and due to wave-particle interactions; and finally, they areinjected into an acceleration process at the termination shockto achieve ACR energies. The predicted abundances are allwithin a factor of 2 of observed values, providing strongvalidation of this scenario. See Schwadron et al. (2002) fora detailed presentation of these points.

5. THE TERMINATION SHOCK FORESHOCK

When V1 reached 84 AU in 2002 it observed increases inenergetic particles with energies from 40 keV to >50 MeV(Stone et al., 2005; Decker et al., 2005). Figure 6 shows thatthese ion fluxes increased in mid-2002 by a factor of 100,disappeared in early 2003, came back at high energies inmid-2003 and at low energies in early 2004, and persisteduntil the termination shock crossing in late 2004. Theseintensity increases were not associated with solar events andwere not observed by V2, which trails V1 by 18.5 AU orabout six years. Thus it seemed likely that these particleswere associated with the termination shock. But these par-ticles were flowing along the magnetic field lines predomi-

Fig. 5. An illustration of ACR production (Schwadron et al., 2002). Lower curves apply to the known interstellar source ACRs (adaptedfrom Jokipii and MacDonald, 1995), while upper curves apply to the outer source, described later in the paper.


nately in one direction and that direction was away fromthe Sun (the actual field direction is tangential due to thewinding of the Parker spiral with distance, so the particleswere moving tangentially but from the direction that is con-nected to the Sun). If the particle source were the termina-tion shock and V1 was inside this shock, then the particlesseemed to be flowing the wrong way. This backward flow

direction led to suggestions that V1 had already crossed thetermination shock (Krimigis et al., 2003) and that the par-ticles were coming outward from the termination shock.Subsequent work has shown that if the termination shockwere blunt, the magnetic field lines first intersect the termi-nation shock closer to the Sun at the distance of V1, so thatparticles from the termination shock would flow to V1 in

Fig. 6. V1 energetic particle data during 2002.0–2005.7 (83.4–97.0 AU). (a), (b), (d) Background-corrected, scan-averaged five-daysmoothed daily-averaged intensities of 40–53-keV ions, 3.4–17.6-MeV protons, and 0.35–1.5-MeV electrons. (c) First-order anisot-ropy vector A1/A0 of proton channel in (b) when the intensity >0.0025 flux units (inset shows orientation of R and T). Whiskersshow direction that particles are traveling. (e) Black bars show times when the V1 plasma wave instrument detected electron plasmaoscillations (Gurnett and Kurth, 2005). (f) Estimated plasma radial flow speed VR based on 40–220-keV ion angular data and a veloc-ity extraction algorithm (Krimigis et al., 2005). Bounds on VR during period A and the second half of period C (2004.56–2004.78) areindicated by the cross-hatched rectangles and means are shown by the horizontal bars. Horizontal bars during 2002.041–2002.172 and2003.257–2003.322 are from Krimigis et al. (2003). Figure courtesy of R. Decker.


the nominally outward direction (Jokipii et al., 2004). Thishypothesis predicts that flows observed by V2, on the op-posite side of the nose of the termination shock, would bein the opposite direction (nominally toward the Sun). In late2004, V2 observed these particles and they were in the op-posite direction from those at V1, as predicted (Decker etal., 2006). Model predictions (Opher et al., 2006) show thatfield lines from the termination shock should intersect V12–3 AU upstream of the termination shock and V2 4–7 AUupstream of the shock. The amount of time each Voyagerspends in the foreshock region depends on the motion of thetermination shock. As V1 approached the termination shockthe solar wind dynamic pressure was increasing, pushingthe termination shock outward, so that V1 was surfing be-hind the shock and remained in or near the foreshock foralmost 9 AU. The solar wind pressure then decreased andthe shock moved quickly inward past V1.

Figure 6 shows that the first termination shock particleevent ended at 2003.1. The end of this event coincided withthe passage of a merged interaction region (MIR) by thespacecraft that pushed the termination shock and thus theforeshock region outward. The first termination shock parti-cle event at V2 also ended when a shock followed by a MIRpassed V2, driving the shock outward. The foreshock par-ticle fluxes observed have a lot of structure and many peaks.Some of this structure is connected with the arrival of MIRsat the location of V1 (Richardson et al., 2005b); since theV1 plasma instrument does not work, the MIR arrival timesat V1 were estimated by propagating V2 data to the dis-tance of V1. These MIRs can affect the particle fluxes inseveral different ways; as in the above example, they canpush the termination shock outward and sever or weakenthe spacecraft connection to the shock. Since the MIRs pushthrough the ambient solar wind, the magnetic field direc-tion may change at the MIR or at the leading shock, thuschanging the magnetic connection between the terminationshock and the spacecraft, which changes the observed parti-cle flux. Since MIRs are regions of enhanced magnetic fieldstrength and turbulence, energetic particles diffuse throughthese regions slowly and pile up ahead of the enhancedmagnetic field region, forming intensity peaks as the par-ticles are “snowplowed” ahead of the MIRs (McDonald etal., 2000). Only one MIR has occurred since V2 entered thesheath; it was associated with an increase of lower-energyparticles (28–540 keV) at the leading shock, probably due toacceleration at the shock. The more-energetic 540–3500 keVparticles had intensity peaks after the shock; these couldresult from snowplowing of particles ahead of the MIR. Theparticle peak could also be behind the shock if the shockhas weakened so that higher-energy particles are no longeraccelerated; the particles already energized before the shocksource turned off would convect with the solar wind whilethe shock would propagate ahead through the solar wind atthe fast mode speed (Decker et al., 2001).

Both V1 and V2 observe periodic enhancements in theparticle fluxes with periods of roughly a solar rotation. Inthe first V1 foreshock encounter and in the V2 foreshock

data through mid-2007, these enhancements are not asso-ciated with changes in the plasma or magnetic field. How-ever, when V1 was near the termination shock, in 2004,many of these enhancements were associated with cross-ings of the heliospheric current sheet (HCS), but only fromnegative to positive magnetic polarities (Richardson et al.,2006b). One explanation for this correlation is that particlesfrom the termination shock current sheet drift (Burger andPotgieter, 1989) along the HCS to V1, and at southern tonorthern magnetic polarity HCS crossings the connectionto the termination shock is much closer. This mechanismwould give peaks once per solar rotation. This hypothesispredicts that when V2, at southerly heliolatitudes, is close tothe termination shock, particle intensity peaks would be ob-served only at positive to negative polarity crossings of theHCS. Particle peaks are observed by both V1 and V2 awayfrom the termination shock with periodicities of roughly asolar rotation; the origin of these peaks is not understood.

6. THE TERMINATION SHOCK CROSSING

The termination shock is where the solar wind plasmamakes its transition from supersonic to subsonic flow. Thesolar wind plasma is compressed and heated and the magneticfield strength increases. Zank (1999) reviews pre-encoun-ter predictions of the nature and location of the terminationshock. Since the average spiral magnetic field upstream ofthe termination shock is along the T axis (where R is out-ward from the Sun and T is in the solar equatorial plane andpositive in the direction of the Sun’s rotation), the termina-tion shock is expected to be a quasi-perpendicular shock.However, energetic particles accelerated or reacceleratedat the termination shock and in the turbulent heliosheathcan propagate into the upstream region. When the Voyagerspacecraft are near the termination shock and connected toit by the magnetic field, measured particle intensities andanisotropies may be comparable to those at the terminationshock if pitch angle scattering is infrequent over the con-nection length and particles can traverse this length beforebeing convected shockward. Otherwise, intensities and an-isotropies observed at Voyager will be reduced relative tothose at the termination shock.

The termination shock was crossed in December 2004 onDOY 351, unfortunately during a tracking gap. The Voy-ager spacecraft are tracked by the DSN 10–12 hours/daydue to competition with other spacecraft. The day before theshock crossing, the V1 Low Energy Charged Particle (LECP)instrument saw an intense beam of field-aligned ions andelectrons at the same time as the Plasma Wave Spectrom-eter (PWS) experiment saw electron plasma oscillations;theory predicts that electron beams will drive these waves(Decker et al., 2005; Gurnett and Kurth, 2005). After theshock crossing the lower-energy (tens of keV) ions jump inintensity by about an order of magnitude above those in theforeshock and the intensities become steady as shown inFig. 6. The termination shock particles shown in Fig. 7 arewell fit by power laws, consistent with diffusive shock ac-


celeration (Decker, 1988) theory for ions energized at ashock. The spectra of termination shock particles both in-side and outside the termination shock are best fit by twopower laws with a break energy at 3.5 MeV. The constancyof the break energy suggests that the termination shock par-ticles in the foreshock and heliosheath have the same source.If the termination shock particles were accelerated by dif-fusive shock acceleration, then the shock strength can bedetermined from the slope of the termination shock particlepower law and is 2.6 (2.4–3.0 with error bars) (Stone et al.,2005). The shock strength is the compression ratio at theshock (Ndownstream/Nupstream) and is 4 for a strong shock. Analternative to diffusive shock acceleration has been sug-gested for the <20-MeV particles; a simple model propa-gating the measured pickup ions from Ulysses outward andheating them at the termination shock can fit the observedspectra at these energies (Gloeckler et al., 2005). In thismodel, up to 80% of the solar wind dynamic pressure goesinto heating the pickup ions. A relatively weak shock is con-

sistent with the magnetic field observations, which showedthe magnetic field strength increased by about a factor of3 at the shock and stayed high, convincing evidence thatV1 was in the heliosheath (Burlaga et al., 2005). The stan-dard deviation of the magnetic field components increasedacross the termination shock, consistent with predictions forthe heliosheath region, and the distribution of the field mag-nitudes changed from log normal in the solar wind toGaussian in the heliosheath. After V1 crossed the termina-tion shock into the heliosheath, the intensities of ions of allenergies first increased, then fluctuated for 20–30 days, andthen became relatively steady, remaining essentially flat atlower energies and increasing slowly at higher energies. Atthe termination shock crossing, the intensity of 40–53-keVions increased tenfold, reached a peak, then dropped andthereafter was relatively smooth, varying typically by nomore than ≈20% about a mean of ≈300 for the next sixmonths. The anisotropies went from beamlike in the fore-shock to nearly isotropic in the heliosheath.

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 termina-tion shock particles is a broken power law with break energy of ~3.5 MeV for protons and a helium. ACR helium dominates the mid-energies in the right panel, with GCRs at higher energies in both panels. The predicted ACR spectra at the shock are shown for astrong (r = 4) and a weak (r = 2.4) shock (Cummings et al., 2002).


The intensity of the more-energetic ACR particles in-creases at the shock but remains lower than the peak valuesin the foreshock. As V1 moved deeper into the heliosheaththe ACR intensities increased. These ions are isotropic, notflowing in field-aligned beams as in the foreshock. The ma-jor surprise at the shock crossing was that the ACR inten-sities did not peak and the ACR spectra did not unroll intoa power law distribution. For almost 30 years the paradigmhad been that ACRs were accelerated at the terminationshock through diffusive shock acceleration. This theory pre-dicted a peak in ACRs at the shock and that the accelerationprocess would result in a power law spectra at the termina-tion shock. Instead, the ACRs intensities at the terminationshock did not reach even the level in the foreshock and theACR energy spectra barely changed. Figure 7 compares theobserved ACR spectra to those predicted for strong andweak shocks. For energies less than 30 MeV the observedintensities are far below those predicted. ACRs were notbeing accelerated where V1 crossed the termination shock.So where are they accelerated? That is still an open ques-tion. The intensities did increase further into the heliosheath.Some investigators have suggested that the heliosheath itselfwas the acceleration region.

Another attractive hypothesis is that the ACRs are accel-erated at the flanks of the heliosphere where the connec-tion time of a field line to the termination shock is longer(McComas and Schwadron, 2006). Field lines first makecontact with a blunt termination shock near its nose. Asthese field lines are dragged out into the heliosheath by thesolar wind, the connection point between the magnetic fieldand termination shock moves toward the flanks. Anomalous

cosmic rays likely require substantial time to be acceler-ated; insufficient connection time near the nose could pre-vent acceleration to ACR energies. The much longer con-nection times toward the distant flanks of the shock mayprovide the necessary acceleration time to achieve the veryhigh 100–300-MeV ACR energies (McComas and Schwad-ron, 2006; Schwadron and McComas, 2007). If the termi-nation shock has a strong nose-to-tail asymmetry, then theflanks would be connected by field lines that are well be-yond V1 near the nose, as illustrated in Fig. 8.

Figure 9 documents the strong correlation between modu-lated ACR and GCR helium. The data points show V1 obser-vations (McDonald et al., 2006) from January 2002 to Febru-ary 2006; the upper panel, energy range 150–380 MeV/nuc, shows GCRs, and the lower panels, energy ranges 10–

Fig. 8. The configuration of the termination shock and magneticfield in the heliosheath from Schwadron and McComas (2007).The bluntness of the termination shock leads to ACR accelera-tion on the distant flanks (McComas and Schwadron, 2006;Schwadron and McComas, 2007). The red region shows the outermodulation boundary of ACRs accelerated on the distant flanksof the termination shock.

Fig. 9. The correlation between modulation of ACRs and GCRs.The panels show the relative fluxes of ACRs and GCRs observedby Voyager 1. The data (plus signs) were reduced by F. McDonaldand B. Heikkila. All fluxes have been normalized by their maxi-mum values and the fluctuations over time (averaged over 26 days)are plotted relative to one another. The respective energy rangesand species are indicated on the axes. The solid curves show theo-retical predictions based on a simple Parker-type diffusion-con-vection modulation model (see text for details).


21 MeV/n and 30–56 MeV/n, measure ACRs. The data arenormalized to the maximum value of the differential energyflux in the respective energy ranges. Schwadron and Mc-Comas (2007) show how the strongly correlated modula-tion of ACRs and GCRs requires that the ACR acceleratorbe on field lines that connect far beyond V1 on either thedistant reaches of termination shock flanks or far out in theinner heliosheath near the heliopause.

Another possibility for explaining ACR acceleration inthe heliosheath is stochastic acceleration (Fisk, 2006). How-ever, stochastic mechanisms encounter several problemswith explaining ACR acceleration at high energies: (1) therates of stochastic acceleration are typically lower than therate of diffusive acceleration at the termination shock, whichmakes it difficult to generate highly energetic ACRs in theone- to a few-year time frame before they become multiplyionized (Jokipii, 1987; Mewaldt et al., 1996); and (2) themechanism must lead to ACRs generated only far out in theheliosheath, or even toward the tail, whereas a distributedprocess seems more likely to generate a very distributedsource. Although stochastic processes may not act quicklyenough to accelerate ACRs to high energies, they are criti-cal for generating the suprathermal seed population, whichis naturally favored for diffusive acceleration at the termi-nation shock (Fisk, 2006).

McDonald et al. (2006) suggest that the lack of an ACRpeak was a time-dependence effect. A MIR passed V1 andencountered the termination shock just before the termina-tion shock crossed V1; perhaps this reduced the accelera-tion efficiency and the ACR spectra recovered with time,so the increase in intensity observed after the V1 termina-tion shock crossing is a time, not radial distance, effect. Thisscenario accounts for the larger ACR fluxes in the foreshockthan at the termination shock; the MIR reduces the ACRfluxes at the termination shock from the previous, largerintensities that were observed flowing into the foreshockfrom the termination shock. Over six months later, the ACRfluxes were still recovering from this large temporal change.This hypothesis will be tested when V2 crosses the termi-nation shock.

Another surprise was that after the termination shockcrossing V1 stayed in the same magnetic sector for 150 days(Burlaga et al., 2005). Since the Sun rotates every 25.4 daysand the HCS is tilted, on average two HCS crossings shouldoccur every 25 days from southern magnetic polarities tonorthern and back. These data were even more puzzlingsince V1 was in the southern polarity sector, where V1should spend less time since it is at northerly latitudes. Sec-tor crossings did resume (Burlaga et al., 2006); the expla-nation for the initial lack of crossings is that the termina-tion shock and heliosheath were moving inward at a speedcomparable to the outward heliosheath flow speed, so V1remained in the same heliosheath plasma as it moved out-ward (Jokipii, 2005). The solar wind pressure was decreas-ing, which would make the termination shock and helio-sheath move inward. The radial speeds deduced from theLECP data range from –30 to 100 km/s but average 30–

50 km/s in this time period, consistent with the terminationshock moving inward.

An interesting feature in the magnetic field data are largefluctuations in the field magnitude, a factor of 3, that arepurely compressional (the field direction does not change)(Burlaga et al., 2006). These features look qualitatively likethe mirror mode waves observed in the magnetosheaths ofJupiter and Saturn (Joy et al., 2006; Bavassano-Catteneoet al., 1998); sometimes they are dips below the backgroundfield value as are observed in low β regions of planetarymagnetosheaths, and sometimes they are increase above thebackground field as observed in high β regions. The periodsof these waves are roughly an hour. If they are mirror modewaves, they result from an instability that occurs in high βplasmas when Tperp > Tpar, where Tperp and Tpar are the tem-peratures perpendicular and parallel to the magnetic fielddirection. These conditions often occur behind perpendicu-lar shocks such as the termination shock. As Voyager getsfurther from the shock, one expects that the plasma will, asa result of these waves, become more stable and the wavesless frequent.

7. THE HELIOSHEATH

The spectra of the termination shock ions in the helio-sheath are much less variable than in the foreshock. Thesteadiness of the spectrum suggests that the spectral shape islittle affected by modulation in the heliosheath, as expected(Jokipii and Giacalone, 2003). The energy spectrum in theheliosheath hardens (has a larger percentage of high-energyions) with time; the intensities of the higher-energy ionsincrease while those of the lower-energy ions are roughlyconstant. The angular distributions of ions from 40 keV toat least 30 MeV were highly anisotropic during the 2.5-yearapproach of V1 to the termination shock. In the heliosheaththe ion anisotropies are much smaller and more variable indirection (Decker et al., 2005; Krimigis et al., 2005).

Figure 6 shows that 5–6 days after the termination shockcrossing, V1 began measuring a positive radial flow com-ponent fluctuating around VR ≈ +100 km s–1, which con-tinued for 27–28 days. Near 2005.045, VR decreased rap-idly (within about 5 days) and flow became inward from2005.063 to 2005.300. From about 2005.30 to 2005.4, VRwas 0 to +30 km s–1. VR increased again around 2005.45and remained near +50 to +100 km s–1 until at least 2005.54.These excursions of VR over a period ≈90–140 days areconsistent with V1 sampling the variable flow downstreamof an inwardly moving termination shock. Convective flowat V1 that fluctuates about zero implies that V1 would es-sentially be riding with and sampling the same parcel ofplasma for an extended period. Tangential flows are also ob-served in the direction away from the heliospheric nose di-rection, as expected for flow around the heliosphere (Deckeret al., 2007).

The relationship between ACRs and termination shockparticles is unknown. Both are deficient in C ions, indicatingthat they are accelerated pickup ions (Krimigis et al., 2003).


However, the H/He ratio is ~10 for termination shock parti-cles (Krimigis et al., 2003) and ~5 for ACRs (Cummingset al., 2002), suggesting that He is more easily acceleratedto ACR energies than H (Zank et al., 2001).

The flow streamlines in the heliosheath were calculatedby Parker (1963), assuming subsonic, incompressible, ir-rotational flow. Using a similar flow model and a kinematicapproximation, Nerney et al. (1991) and Washimi and Ta-naka (1996) calculated the variation of the magnetic fieldB in the heliosheath. Close to the termination shock, the ve-locity remains nearly radial and B remains nearly azimuthal.The speed decreases as the plasma moves toward the helio-pause; consequently, the magnetic field strength B in-creases. Cranfill (1971) and Axford (1972) showed that ifthe heliosheath is thick enough, B might become strongenough to influence the flow. Nerney et al. (1993) estimatedthat B might influence the flow in a restricted region. Theflow is deflected away from the radial direction as it ap-proaches the nose of the heliopause, and B develops a radialcomponent. At the heliopause, B must be parallel to the sur-face (unless significant reconnection occurs at the bound-ary). The observations of B in the heliosheath show fea-tures not described by the simple kinematic models. WhileB does have a strong azimuthal component, a persistent me-ridional component (northward) is observed that was notpredicted (Burlaga et al., 2005).

The magnetic field direction in the solar wind is oftenalternately toward and away from the Sun for several days.The existence of these “sectors” is related to extensions offields from the polar regions of the Sun to the latitude of theobserving spacecraft. Sectors have been observed by V1 outto 94 AU (Burlaga et al., 2003b, 2005).

8. THE HELIOPAUSE

As the Voyagers continue their trek through the un-charted heliosheath, the next milepost in their journey willbe the heliopause, which we think they will encounter be-tween 2013 and 2021. The crossing of the heliopause willbe a truly historic occasion; the first manmade vehicle toleave the solar system. The heliopause is thought to be atangential discontinuity with a large jump in the plasmadensity and a rotation and change in the magnitude of themagnetic field. Given the lack of an operational plasma in-strument on V1, the first crossing of the heliopause willprobably be identified by a durable change in magnitudeand/or direction of the magnetic field. If the model asym-metries discussed above are real, then the V1 and V2 cross-ings could be close to each other in time.

The flow velocity vectors are expected to become tan-gent to the nominal heliopause surface. However, variationsof the heliopause surface and fluctuations in the flow ve-locities could make it difficult to use these data to identifya heliopause crossing. In analogy with planetary magneto-pauses, the heliopause is probably a complex surface thatvaries locally in thickness and orientation and is the site ofpatchy reconnection and a variety of plasma instabilities.

The appearance of large variations in flow velocities associ-ated with disturbances from such relatively small-scale dy-namical processes may indicate that V1 is near the helio-pause. Reconnection at the heliopause-LISM interface couldaccelerate low-energy ions (and electrons) and create largedepartures from the normal heliosheath intensities andanisotropies. Intermittent reconnection along a slightly cor-rugated heliopause surface might appear in the low-energyion (and possibly, electron) intensity-time profiles as grad-ual intensity increases with superposed, anisotropic inten-sity spikes as the spacecraft approached the heliopause, notunlike structures encountered by V1 as it approached thetermination shock. These comments are speculative, draw-ing largely upon analogies to planetary magnetopauses;however, as with the termination shock, energetic particleobservations may enable us to remotely sense the helio-pause well before the Voyagers cross it.

Instabilities on the heliopause have been the subject oftheoretical conjecture (cf. Fahr, 1986; Florinski et al.,2005). In principle, the Rayleigh-Taylor instability could beimportant near the nose of the heliosphere where the dif-ference between the flow densities are large. This instabilityrequires an effective “gravity” or destabilizing force. Lieweret al. (1996) proposed that coupling between ions and neu-trals via charge exchange would produce such a destabiliz-ing force. Near the flanks, away from the nose, the Kelvin-Helmholtz instability is more likely to be important wherethe shear velocity across the heliopause is large. Either typeof instability would produce motion of the heliopause su-perposed on the “breathing” expected to occur due to vary-ing solar wind pressure. Estimates of the magnitude of theresulting excursions are tens of AU and have periods on theorder of 100 years or more (Liewer et al., 1996; Wang andBelcher, 1998). For Voyager, the long periods make it un-likely that such wave motion can be detected, but the addi-tional radial motion from such instabilities, should theyexist, increases the uncertainty in the distance to the helio-pause.

9. THE HYDROGEN WALL

Plate 11 shows that a stagnation point forms in the flowat the nose of the heliosphere. The plasma flow speed inthis 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 Baranovand Malama, 1993, 1995; Zank et al., 1996). Plate 11 showsthe region of denser H known as the hydrogen wall in frontof the heliosphere. This wall has been observed. Observa-tions of H Lyman-α emission from nearby stars in the nosedirection show that the H absorption is red-shifted relativeto the other lines by 2.2 km/s; its width is too broad to beconsistent with the T = 5400 ± 500 K temperature suggestedby the width of the deuterium line (Linsky and Wood, 1996).The profiles could only be fit assuming additional absorp-tion by hot H, such as the compressed H in the hydrogen


wall. Two H components were needed to fit the data, onefrom the heliospheric hydrogen wall (which, since the H ismoving sunward slower than the local interstellar cloud H,produces absorption on the red side of the H line) and onefrom the hydrogen wall of the nearby stars (which, sincethe H is moving sunward faster than the local interstellarcloud H, produces absorption on the blue side of the H line)(Linsky and Wood, 1996; Gayley et al., 1997; Izmodenovet al., 1999, 2002; Wood et al., 1996, 2000). Many hydro-gen walls have been detected in front of other astrospheres,allowing comparison to the heliosphere and providing vali-dation to our general view of the heliosphere structure.

10. THE HELIOSPHERIC RADIO EMISSIONS

The Voyager spacecraft periodically detect radio emis-sion from the boundaries of our solar system. This radiosource is the most powerful in the solar system with anemitted power of over 1013 W (Gurnett and Kurth, 1996).Plate 12 shows a spectrogram of these emissions, plottingintensity as a function of time and frequency. These waveshave two frequency components, a constant frequency 2-kHz component and a component where the frequency risesfrom 2.4 to 3.5 kHz over a period of 180 days. The twolargest events occurred in 1983–1984 and 1992–1994, a fewyears after solar maxima. Large interplanetary shocks orMIRs were hypothesized to trigger these emissions whenthey entered the higher-density hydrogen wall beyond theheliopause. But despite the occurrence of several large eventsafter the 2000 solar maximum, only very weak emissionwas observed in 2002–2003, even though Voyager is pre-sumably much closer to the radio source. The source re-gion and trigger for these emissions is not well understood.

11. THE HELIOSHEATH: THE LOCALINTERSTELLAR MEDIUM

The local interstellar medium surrounding the helio-sphere is not a constant. As the Sun and interstellar cloudsmove through the galaxy, the Sun has been and will be invery different environments. The interstellar medium hashot tenuous regions and cool dense regions that coexist inpressure balance. The Sun is now in a hot tenuous regionof the interstellar medium known as the local bubble. Thisregion is about 200 pc across and has T = 106 K and anaverage neutral H density of 0.07 cm–3. This density is about10 times smaller than the average interstellar H density. Thelocal bubble was created by several supernova explosionsthat occurred on the order of a million years ago, blowinga hole in the interstellar medium. The local bubble containsseveral clouds, including the local interstellar cloud in whichthe Sun now resides, with an H density of about 0.2 cm–3.The Sun is very near the edge of the local interstellar cloudand will enter a different interstellar environment in less than4000 years. It entered the current local interstellar cloudwithin the past 40,000 years. Possible values of the densitiesencountered by the heliosphere in the past and future are0.005–15 cm–3 with Sun-cloud relative velocities of up to

100 km/s (Müller et al., 2006). Given possible values of theinterstellar pressure, in the past and future the heliopausecould be located anywhere between 12 and 400 AU and thetermination shock from between 11 and 250 AU. For a verysmall heliosphere, cosmic-ray fluxes would dramatically in-crease, affecting the climate and increasing rates of geneticmutation. When the heliosphere is in the local bubble, itwill be greatly expanded. The GCR fluxes at Earth will besimilar to present-day fluxes, but the fluxes of ACRSand neutral H will be greatly reduced.

12. FUTURE OBSERVATIONS

12.1. Voyager

The Voyager spacecraft continue their journeys, V1 mov-ing out about 3.5 AU/year and V2 about 3 AU/year. Voy-ager 2 is in the termination shock foreshock region; mod-els suggest that this region only extends 4–7 AU from thetermination shock. In June 2007, V2 had already been inthe foreshock region for almost 8 AU. The location of thetermination shock and foreshock change with the solar windpressure and thus the time spent in the foreshock can belonger than its width would suggest, as when the termina-tion shock was moving out just ahead of V1, but V2 willprobably cross the termination shock before the end of 2008.This crossing will provide information on the asymmetryof the heliosphere. It will also provide the first plasma datafrom the heliosheath, providing direct information on theplasma flow, shock compression, and heating at the termi-nation shock. Figure 10 shows the V1 and V2 trajectoriesin distance and heliolatitude. Model estimates of the loca-tions of the termination shock and heliopause are shownshown by the lines.

The V1 heliopause crossing is likely to occur between2013 and 2021 and the V2 heliopause crossing should occurin the same time frame. The spacecraft will have sufficient

Fig. 10. Trajectories of the Voyager spacecraft in distance andheliolatitude. The range of predicted heliopause (for V1 and V2)and termination shock (for V2) distances are shown by the solidparts of the trajectory.


power to operate all instruments until 2016; after this time,power-sharing will extend the useful life of the spacecraftbeyond 2020. Thus the Voyagers are likely to provide thefirst in situ measurements of the local interstellar cloud.

12.2. New Horizons

New Horizons is a NASA mission that will fly by thePluto-Charon system and then proceed to other Kuiper beltobjects. Again, due purely to serendipity, Pluto is now lo-cated toward the nose of the heliosphere. Although thismission will not return data from the interstellar medium,it will measure the pickup ions directly, providing new in-formation on the densities and composition of the interstel-lar neutrals entering the solar system.

12.3. Interstellar Boundary Explorer

The Interstellar Boundary Explorer (IBEX) is an excit-ing new NASA mission that will observe the large-scalestructure of the heliosphere directly. It will be launched intoa highly eccentric Earth orbit in June 2008 [see the gen-eral description of the mission in McComas et al. (2004)].All-sky energetic neutral atom (ENA) images of the globalstructure of the termination shock and heliosheath will beaccumulated every 6 months. These ENAs are generated bycharge-exchange collisions between the inflowing interstel-lar neutral gas and the protons that are heated/acceleratedat the termination shock and that subsequently populate theentire heliosheath. Two IBEX single-pixel telescopes willimage these hydrogen ENAs over semiannually precess-ing great circles in the sky at low energies (10 eV to 2 keV)and in an overlapping band of higher energies (300 keV to6 keV). The intensities of the ENAs depend strongly on boththe shape of the proton spectrum and the plasma convec-tion flow pattern in the heliosheath. Where the anti-Sunwardcomponent of the heliosheath convection velocity is high,the ENA intensity directed toward Earth is low, and viceversa. The ENA foreground from the heliosphere betweenEarth and the termination shock is negligible because thesolar wind and pickup ion populations are being convectedaway from the Sun at the solar wind velocity (which isgreater than the Earth-directed velocities of the thermal orpickup ions). Beyond the termination shock, the shockeddecelerated flow in the heliosheath allows heated/acceler-ated protons to reach Earth and be imaged by IBEX.

The ENA emission seen from Earth is accumulated alongthe entire line of sight outward through the inner helio-sheath, across the heliopause, and into the outer heliosheath.The measured ENA intensity integrates the effect of the ter-mination shock heating/acceleration and heliosheath plasmaflow over several tens of AU, thus yielding diagnostics ofthis immense boundary region (Gruntman et al., 2001).Even though V1 and V2 do not directly measure protonsin the energy range corresponding to the IBEX ENA hy-drogen, these complementary measurements will prove ofinestimable value in constraining the theoretical models ofthe global interaction of the solar wind with the interstel-

lar medium. For example, Gloeckler et al. (2005) extrapo-lated the V1 LECP proton spectra (40–4000 keV) upstreamand downstream of the termination shock to lower energiesand demonstrated that 80% of the upstream ram pressuregoes into heating the pickup protons. Thus V1 providesthe input proton spectrum for the global heliosheath (albeitonly at the V1 location) that will be imaged in IBEX hy-drogen ENAs. And, returning the scientific favor, the veryfirst IBEX half-year images will, at a glance, reveal asym-metries in the large-scale heliosheath configuration that bearon the puzzling directions of the magnetic field and ener-getic particle streaming observed by V1 and V2. The IBEXall-sky images, in concert with local measurements fromVoyager, will provide a quantum leap forward in our un-derstating of the outer heliosphere and its interstellar inter-action.

12.4. The Interstellar Probe

Finally, in the era after Voyager and IBEX, we hope tosend a new, optimally instrumented mission called the In-terstellar Probe out beyond the termination shock, helio-pause, and bow shock and into the pristine interstellar me-dium. Such a mission will require propulsion capable ofaccelerating a spacecraft to immense speeds if results areto be gained in a reasonable mission lifetime. Recent studieshave examined both solar sails and nuclear-powered elec-tric propulsion. While the results of such a mission are stilldecades off, we need to begin soon if we are to finally senda probe out into the material of the stars.

Acknowledgments. J.D.R. was supported under NASA con-tract 959203 from the Jet Propulsion Laboratory to the Massachu-setts Institute of Technology. N.A.S. was supported by the NASAInterstellar Boundary Explorer mission, which is part of the GSFCExplorer Program, and by NASA’s Earth-Moon-Mars RadiationEnvironment Model (EMMREM) Project.

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