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Home , Bow shock, Heliosphere, Lambda Andromedae, Local Bubble, Local Interstellar Cloud, Pickup ion

Richardson and Schwadron: The Limits of Our 443

The Limits of Our Solar System

John D. Richardson Massachusetts Institute of Technology

Nathan A. Schwadron Boston University

The is the bubble the carves out of the . The solar moves outward at supersonic speeds and carries with it the Sun’s . The interstellar medium is also moving at supersonic speeds and carries with it the interstellar magnetic field. The boundary between the 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. , which recently crossed the termination shock of the solar wind, is the first 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 heliopause is the inner heliosheath and the region be- tween the and the heliopause is the outer helio- The limits of the solar system can be defined in many sheath. ways. The farthest object orbiting the Sun could be one. The Plate 11 shows the equatorial plane from a model simu- farthest could be another. We define this limit as the lation of the heliospheric system (Müller et al., 2006). The boundary between the Sun’s solar wind and the interstellar color coding in the top panel gives the temperature medium, the material between the in our . The and in the bottom panel shows the neutral H density. The interstellar medium is not uniform, but is composed of many plasma flow lines are shown in the top panel. The trajecto- clouds with different densities, temperatures, magnetic field ries of Voyager 1 (V1) and (V2) are also shown; strengths and directions, and flow speeds and directions. V1 is headed nearer the heliospheric nose than V2. We note The Sun is now embedded in a region of the interstellar me- that it is only by serendipity that the Voyager spacecraft are dium called the local . The solar wind flows headed toward the nose of the heliosphere; the Voyager tra- outward from the Sun and creates a bubble in the interstel- jectories were driven solely by the locations of the lar medium. that were part of Voyager’s grand tour of the solar system. The solar wind pushes into the The main heliospheric boundaries, termination shock, helio- until it too weak to push any farther, then turns and moves pause, and bow shock are labeled in the upper panel. The downstream in the direction of the local interstellar cloud local interstellar cloud is only barely supersonic, so the heat- flow. The plasmas in these two cannot mix, since they ing that occurs at the bow shock is small; at this boundary are confined to magnetic field lines and magnetic fields can- the flow of local interstellar cloud plasma begins to divert not move through one another. Neutrals, however, are not around the heliosphere. In the solar wind the flow is radi- bound by magnetic fields and can cross between regions. So ally outward until the termination shock is reached. At the local interstellar cloud neutrals can and do penetrate deep termination shock, the plasma is strongly heated and the into the solar system. The boundary between the solar and flow lines bend toward the tail. At the heliopause, the local interstellar winds is called the heliopause. interstellar cloud and solar wind flow are parallel; plasma Both the solar wind and the interstellar wind are super- generally cannot cross this boundary. The neutral density sonic; thus a shock forms upstream of the heliopause in each increases at the nose of the bow shock, forming an enhanced flow. At these shocks, the plasmas in each wind are com- density region known as the wall, and decreases pressed, heated, become subsonic, and change flow direc- at the heliopause. tion so that the plasma can move around the heliopause (for We are just beginning to learn about these outer regions interstellar medium) or down the heliotail (for the solar of our heliosphere and how they interact with the local in- wind). The shock in the solar wind is called the termination terstellar cloud. Until V1 crossed the termination shock in shock. The shock in the local interstellar cloud is called the 2004 at 94 astronomical units (AU) (1 AU is the distance bow shock. The region between the termination shock and from the Sun to ), we could only guess the boundary

443 444 The Solar System Beyond locations. Now we have observed one termination shock spirals become tightly wound and the magnetic field angle crossing and have several of heliosheath data, but we is nearly 90° (the field is almost purely tangential). Since still have much to learn about the other heliosphere bounda- the magnetic field lines from the northern and the southern ries and the global aspects of these boundaries. This chapter hemispheres of the Sun have opposite polarity, a helio- discusses what we know about the interaction of the helio- spheric current sheet (HCS) forms near the equator, across sphere with the local interstellar cloud and these boundary which the magnetic field direction reverses. The HCS passes regions, the unsolved problems, and the future observations through the termination shock and has been observed in the we hope will help answer these questions. 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, 2. THE HELIOSPHERE an average temperature of 96,000 K, and an average mag- netic field strength of about 5 nT; these values are based The heliosphere is the region dominated by plasma and on a of data from the Wind spacecraft from 1994 magnetic field from the Sun. No spacecraft has yet crossed to 2005. The average dynamic is about 2.25 nP; outside this region; estimates of the size of the heliosphere this pressure determines how big a cavity the solar wind at the nose (in the upwind direction) range from 105 to carves into the local interstellar cloud. 150 AU (Opher et al., 2006). In the tail, the heliosphere extends many hundreds and probably thousands of AU. The 2.2. The Local Interstellar Cloud location of the nose of the heliopause is determined by pres- sure balance between the outward dynamic pressure of the The local interstellar cloud surrounds the heliosphere. It solar wind and the pressure of the local interstellar cloud, contains neutral and ionized low-energy particles, H, He, O, which has contributions from the dynamic pressure, thermal H+, He+, and O+, as well as very-high-energy galactic cosmic pressure, and magnetic pressure (see Belcher et al., 1993). rays (GCRs). These GCRs are observed at Earth but their The dynamic are given by 1/2ρV2 where ρ is the intensities are modulated (reduced) by the solar wind mag- mass density and V is the speed. For the local interstellar netic field as they move through the heliosphere. Plate 11 cloud the magnetic pressure B2/8π and the thermal pressures shows how the low-energy plasma deflects around the helio- nkT could be important but the magnitudes are not well sphere. Helium is the only major species that makes it to constrained. We discuss first the unperturbed flows from the Earth in pristine form and thus provides the best determi- Sun and in the local interstellar cloud, then discuss the in- nation of the properties of the interstellar medium. Obser- teractions of these two flows. vations of the H and He emission combined with models of the neutral flow give estimates of the neutral speed, den- 2.1. The Solar Wind sity, and temperature as well as the flow direction. Hydrogen is the largest component of the interstellar me- The Sun is a source of both photons and plasma; the dium, but it is difficult to determine the local interstellar fluxes of each fall off as 1/R2. The effects of solar cloud properties from H observations because a large frac- on the heliosphere/local interstellar cloud interaction are tion of the H charge exchanges with the local interstellar small compared to those of the solar wind. The first detec- cloud H+ before it enters the heliosphere. Charge exchange tions of the solar wind in the early 1960s (Neugebauer and occurs when an ion and a neutral collide and an electron Snyder, 1962; Gringauz, 1961) confirmed Parker’s (1958) goes from the neutral to the ion. The former ion is a now theoretical predictions of the basic solar wind structure. The neutral and is not bound by magnetic fields, so it continues solar wind has been thoroughly measured and monitored to move at the speed it had when it was an ion. The former by spacecraft since that time. In addition to the near-equa- neutral, now an ion, is bound to the magnetic field and torial spacecraft in the inner heliosphere, Pioneers 10 and accelerated to the speed of the bulk plasma flow. Outside 11 and V1 and V2 have observed the solar wind in the outer the bow shock the plasma and neutrals have the same speed heliosphere, in the case of V1 past 100 AU. has pro- and temperature so charge exchange has no effect. But in- vided measurements in the high-latitude heliosphere. These side the bow shock, the plasma parameters are different observations give a good picture of the solar wind parame- from those in the interstellar medium since the plasma is ters throughout the heliosphere. heated, compressed, and deflected at the shock. The plasma The solar wind is a supersonic flow of mostly protons, slows further as it piles up in front of the heliosphere. So with on average 3–4% He++ (this percentage varies over a when charge exchange occurs, the new neutrals have speeds solar cycle) (Aellig et al., 2001) and a small amount of and temperatures from the shocked plasma, not the pristine highly charged heavy ions. The Sun has a magnetic field interstellar medium. The H observed in the heliosphere has that reverses polarity every 11- solar cycle. To first or- an inflow speed of only 20 km/s, 6 km/s less than the inter- der, plasma cannot cross magnetic field lines, so the solar stellar medium speed (Lallement et al., 1993), confirming wind outflow drags the solar magnetic field lines outward. that the H has charge exchanged with local interstellar cloud This outward motion combined with the Sun’s 27-day rota- H+, which has slowed down near the heliosphere boundary. tion causes the magnetic field lines to form spirals; at Earth The plasma flow in this region has begun to divert around the spiral angle is about 45°. In the outer heliosphere these the heliosphere, so the flow direction of the H is also altered. Richardson and Schwadron: The Limits of Our Solar System 445

Observations with the SOHO/SWAN instrument show that The gas detector on Ulysses measures the neutral He the arrival direction of the interstellar H differs from that velocity and temperature directly (Witte, 2004; Witte et al., of the interstellar He by about 4° (Lallement et al., 2005). 1992). Even with these direct measurements, determination This offset results from charge exchange of the H with the of the density of the local interstellar cloud still requires deflected H+ in front of the heliosphere. The direction of modeling since some of the He is lost on its way through the the offset tells us the direction of the plasma flow in the heliosphere. All these measurements taken together provide local interstellar cloud just outside the heliopause and thus fairly tight constraints on the local interstellar cloud He neu- provides information on the magnetic field direction (Lalle- tral speed and temperature (Möbius et al., 2004). ment et al., 2005). The magnetic field strength in the local The He moves toward the Sun from an ecliptic longitude interstellar cloud is not known; generally it is thought to of 75° and an ecliptic latitude of –5° with a speed of 26.4 ± be on the order of 1 µG with an upper limit of 3–4 µG. 0.3 km/s. The temperature is 6300 ± 340 K and the density The interstellar neutrals enter the heliosphere and absorb of He is 0.015 ± 0.002 cm–3. The collision timescales in the and emit at specific wavelengths. Backscattered H local interstellar cloud are less than the neutral and plasma Lyman-α (Bertaux and Blamont, 1971; Thomas and Krassa, lifetimes, so the plasma and neutral species should all have 1971) and He 58.4 nM radiation have both been observed the same speed and temperature far from the heliosphere. (Weller and Meier, 1979). As the neutrals move toward the Since the H is altered by charge exchange with the plasma Sun, they are accelerated inward by gravity but pushed out- near the heliosphere, the H density is more difficult to de- ward by the radiation pressure (see review by Fahr, 1974). termine. Best estimates are on the order of 0.20 ± 0.03 cm–3 For H, the radiation pressure and gravity are comparable, in the local interstellar cloud. Since the ionization fraction so the H is not accelerated inward. Helium is heavier and depends on the electron temperature, which is not well is accelerated inward. Hydrogen also has a much shorter known, the plasma density is also uncertain, with best esti- ionization time in the solar wind than He. The result is that mates of 0.06–0.10 cm–3 based on the speed of the H meas- most H becomes ionized before it reaches 5–10 AU. The ured at Earth (Lallement et al., 1993). peak in backscattered Lyman-α intensity is roughly in the direction of the local interstellar cloud flow with a small 3. THE TERMINATION SHOCK LOCATION deflection occurring from the charge exchange in the hy- drogen wall. Although some properties of the local interstellar cloud The ionization time of He is long; most He is not ionized are well constrained, others are not, which makes determin- until it is inside 1 AU. The He with trajectories that remain ing the heliosphere size and shape difficult. We have ob- outside 1 AU survive and flow downstream out of the helio- served one boundary location; the termination shock at the sphere. The Sun’s gravity bends the trajectories so that the position of V1 was at 94 AU on December 16, 2004 (Stone He is focused downstream of the Sun, where the maximum et al., 2005; Decker et al., 2005; Burlaga et al., 2005). This emission from He is observed (Axford, 1972; Fahr and Lay, crossing provides one point with which to calibrate our 1973). The location of the maximum He emission is thus models. Since the average solar wind dynamic pressure is opposite to the direction of the incoming He flow. Obser- 2.25 nP at 1 AU, one can roughly estimate that the total vation of the direction with peak He intensity is one method local interstellar cloud pressure must be 2.5 × 10–4 nP. But used to determine the local interstellar cloud flow direction. there are problems with this simple approach, which we de- The interstellar neutrals are ionized in the solar wind via scribe below. charge exchange with protons (the dominate process for H) or photoionization. These neutrals can be observed directly 3.1. Solar Wind Time Dependence as H+ and He+ pickup ions. He+ is rarely observed in the solar wind and thus is a good marker of interstellar neutrals. The average solar wind values given above are only Both H+ and He+ pickup ions have distinctive signatures in a piece of the story. The solar wind pressure is variable the solar wind; when the neutral H and He become ionized on scales from hours to 11-year solar cycles and likely has they are accelerated to the solar wind speed. These ions, longer-term variations. Small-scale variations of the pres- called pickup ions because they are picked up (accelerated) sure do not affect the boundary locations, but larger-scale by the solar wind, have a gyromotion roughly equal to the variations change the balance between the solar wind and solar wind speed, so these ions show a sharp energy cutoff local interstellar cloud pressures and cause the termination at four times the solar wind energy. shock and heliopause to move in and out. The solar wind The energy to accelerate the pickup ions comes from the dynamic pressure changes by a factor of about 2 over a solar bulk flow of the solar wind, so an indirect way to measure cycle, with minimum pressure near (when the density (discussed further below) is to look at numbers are highest), then a rise in pressure with how much the solar wind slows down as it moves outward. a peak a few years after solar maximum, followed by a de- These direct and indirect observations can be combined with crease to the next solar maximum (Lazarus and McNutt, models of the neutral inflow and loss to estimate speeds, 1990). Models of the heliosphere show that these pressure temperatures, and densities of H and He in the interstellar changes drive an inward and outward “breathing” of the he- medium. liosphere; the termination shock location changes by about 446 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 , the solar wind near the equator is slow (400 km/ s) and dense (8 cm–3). The solar wind at higher latitudes comes from 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 low- speed solar wind. Despite these latitudinal changes in the solar 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 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. Fig. 1. The density of the plasma and neutral components from When a series of ICMEs occurs, the later ones catch up 1 to 1000 AU. The solar wind ions come from the Sun. The pickup to and merge with previous ICMEs, causing a buildup of ions are interstellar neutrals that have been ionized in the solar plasma and magnetic field known as a merged interaction wind. The densities of the solar wind and of the pickup ions jump region (MIR) (Burlaga et al., 1984; Burlaga, 1995). These at the termination shock. Outside the heliopause, the ions are part regions are usually preceded by a fast forward shock at of the local interstellar cloud. Both the ion and neutral density increase in front of the heliopause (in the so-called hydrogen wall). which the speed, density, temperature, and magnetic field The interstellar neutrals dominate the mass density outside 10 AU. strength all increase. Shocks are rare in the outer helio- Figure courtesy of R. Mewaldt. sphere; most can be related to a large CME or a series of CMEs on the Sun (see Richardson et al., 2005a). The dy- namic pressure in the MIRs can be a factor of 10 higher than that in the ambient solar wind and these structures last 10 AU. Figure 1 shows the density of plasma and neutrals 1–3 months (Richardson et al., 2003). The MIRs drive the vs. distance (we know now the termination shock and he- termination shock outward several AU for 2–3 months, after liopause in this figure are located too close to the Sun). The which the termination shock rebounds inward (Zank and solar wind plasma falls off as R–2, then increases abruptly Müller, 2003). at the termination shock and heliopause. The H density is about 0.1 at the termination shock and 0.2 in the local in- 3.2. Local Interstellar Cloud Time Dependence terstellar cloud, then decreases rapidly inside 10 AU, the ionization cavity inside which most H is ionized. Both the Changes in the pressure or magnetic field of the local H and H+ increase between the bow shock and heliopause interstellar cloud could also cause the boundary locations to and form the hydrogen wall. The pickup ion density/thermal move. The scale length for changes in the local interstellar ion density ratio increases to the termination shock, where cloud is not known. The Sun is very near the local inter- about 20% of the ions are pickup ions. This ratio is consis- stellar cloud boundary and will cross into another interstel- tent with the observed slowdown of the solar wind. Outside lar medium environment in a few thousand years, and scale about 30 AU the pickup ions (of local interstellar cloud ori- lengths for variations may be smaller near this boundary gin) dominate the plasma thermal pressure. So even though (Müller et al., 2006). Observations of He by Ulysses are the Sun has carved out its own space in the local interstel- consistent with no change in the interstellar neutral density lar cloud, the local interstellar cloud influence is very strong. since the launch of Ulysses in 1990 (Witte, 2004). 3.3.1. Solar wind slowdown. The neutral interstellar H is continuously being lost, mostly by charge exchange with 3.3. Local Interstellar Cloud Effects solar wind protons. The neutrals formed from the solar wind on the Solar Wind protons escape the solar system at the solar wind speed. The new protons formed from ionized local interstellar cloud If we assume that the local interstellar cloud pressure neutrals are accelerated from their initial speed of 20 km/s were constant, we can use the observed solar wind pres- inward to the solar wind speed, 400–700 km/s outward. As sures to model the heliosphere boundary changes. Models mentioned earlier, plasma does not flow across field lines; must include the effects of the local interstellar cloud neu- ions at speeds different from the bulk solar wind speed are trals on the solar wind. These neutrals are the largest (by accelerated by an electric field E = –VXB. The initial ther- mass) component of the interstellar space outside about mal energy of the pickup ions is equal to the solar wind en- Richardson and Schwadron: The Limits of Our Solar System 447

ergy, about 1 keV. The energy for the acceleration and heat- nation shock with an uncertainty of 0.01 cm–3. We note that ing comes from the bulk flow of the solar wind. the H density at the termination shock is not the density in Two consequences of these pickup ions are that the so- the local interstellar cloud; much filtration, or H loss, oc- lar wind is heated and slows down. The amount of the slow- curs between the bow shock and the termination shock. down depends on the density of the local interstellar cloud By 95 AU, the solar wind speed has decreased by about H, so measurements of the slowdown give an estimate of the 20% [and since the mass flux, the speed (v) times the mass local interstellar cloud density. This slowdown can be dif- density (ρ), is a constant, the density must increase by 20%]. ficult to measure because the solar wind speed varies with So the dynamic pressure, ρv2, is only 77% of what it was at time and solar latitude, as well as with distance. In the in- 1 AU. The effect of the interstellar H is to lower the solar ner heliosphere, the slowdown is small and difficult to de- wind dynamic pressure. This reduced pressure causes the tect. In the outer heliosphere, determining the slowdown heliosphere boundaries to be closer to the Sun. The solar from the inner heliosphere to V2 requires (1) another space- wind slowdown is related to the pickup ion density by npu/ γ γ δ craft in the inner heliosphere at a similar heliolatitude to nsw = (3 – 1)/(2 – 1)( v/v0), where nsw is the solar wind δ provide a speed baseline and (2) a model to propagate the density, npu is the pickup ion density, v is the solar wind solar wind outward to determine how much of the speed slowdown, v0 is the solar wind speed at 1 AU, and is the γ δ change comes from radial evolution of the solar wind and ratio of specific heats. For = 5/3, npu/nsw = 7 v/6v0 how much is due to pickup ions. Figure 2 shows a compi- (Richardson et al., 1995). lation of various determinations of the slowdown. The plot 3.3.2. Solar wind heating. As the solar wind moves shows the solar wind speed normalized to 1 AU. The hori- out, it expands and thus might be expected to cool. The solar zontal lines show the average amount of slowdown; the wind temperature does decrease out to about 30 AU, but length of these lines shows the radial range covered in each not as quickly as adiabatic cooling would predict, then in- study. The vertical spread shows the uncertainties in the creases with distance with some solar cycle variations su- slowdown determination. The curves shows the predicted perposed. Internally driven processes, such as shocks and solar wind slowdown for densities of interstellar H at the dissipation of magnetic fluctuations, provide some heating, termination shock of 0.07, 0.09, and 0.11 cm–3. The data but not enough to produce the observed temperature pro- are best fit with a density of about 0.09 cm–3 at the termi- file (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 trans- fer 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. Termination Shock Model Results

Many models developed to determine the termination shock 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 that includes the effect of the local interstellar cloud neutrals on Fig. 2. The plot shows the slowdown of the solar wind caused the solar wind and use it to track the termination shock loca- by the interstellar neutrals. The fraction of the initial solar wind tion. Figure 3 shows predictions of the termination shock speed is plotted vs. distance. The horizontal lines with error bars location based on a two-dimensional model that includes the show observed speed decreases; the lengths of these lines show the radial width over which the slowdown was observed. The effect of the local interstellar cloud neutrals on the solar vertical lines are error bars. The curves shows model predictions wind. The model input is the solar wind pressure observed of the solar wind speed as a function of distance for interstellar by V2; the profile is normalized to the V1 termination shock neutral H densities of 0.09 cm–3 (solid line), 0.07 cm–3 (upper crossing (Richardson et al., 2006a). The model shows that dashed line), and 0.11 cm–3 (lower dashed line). V1 and the shock were both moving radially outward start- 448 The Solar System Beyond Neptune

gential to the shock that excites these emissions. In this case, the magnetic field would be roughly perpendicular to the . 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 helio- sphere 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 inter- . 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 asym- Fig. 3. The predicted position of the termination shock using V2 metry in the model heliosphere in both the north/south and observations of the solar wind as input to a two-dimensional time- east/west directions. The models show that magnetic field dependent hydrodynamic model. The termination shock locations lines that intersect the termination shock extend 2 AU in- are extended into the future by repeating the pressure profile from side the termination shock at the position of V1 and 4–7 AU the previous solar cycle. The V1 and V2 trajectories are shown inside the termination shock in the direction of V2 (Opher et by the dashed diagonal lines. The lower curve shows the termi- nation shock position shifted inward 10 AU to simulate a possible al., 2006). These distances represent the average thickness termination shock asymmetry. The termination shock moved in- of the foreshock region in front of the termination shock. ward from 1995 to 1999, then moved outward ahead of V1 until Small-scale magnetic field variations will affect these dis- 2004, when it moved inward and crossed V1. The model suggests tances locally. This model also predicts that the distance of V2 could cross the termination shock before the end of 2007. the termination shock in the direction of V2 is 8–11 AU closer than in the direction of V1 and the heliopause dis- tance is 13–24 AU closer in the direction of V2. This model did not include neutrals, which may reduce the asymmetry ing in 1999, that V1 made some close encounters to the (Pogorolev et al., 2006). Other models also give asymme- shock, but that the shock crossing did not occur until the tries with ratios of the termination shock distances at V1 termination shock moved inward in 2004. 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 3.5. Heliospheric Asymmetries V2, with heliosheath thicknesses of 55–67 AU toward V1 and 33–50 AU toward V2 (see Opher et al., 2006). At Earth, the 45° average angle of the solar magnetic None of these models include all the important ; field to the solar wind flow direction results in asymmetries that task is currently too expensive computationally. In par- in the plasma (Paularena et al., 2001) and ticular, the effects of the HCS may be important, since the the locations (Dmitriev et al., 2004). The location of the HCS at the termination shock varies over a draping of magnetic field lines around the solar rotation, but this small-scale structure is hard to in- results in different field strengths on the dawn and dusk clude in a global model. Although the models differ quan- sides, causing these asymmetries. The interstellar magnetic titatively, the results suggest that the average shock distance field could result in a similar distortions of the termination at V2 could be ~5–10 AU closer than at V1. In addition, shock and heliopause and significantly affect the width of the solar cycle variation of the solar wind dynamic pressure the heliosheath (see, e.g., Linde et al., 1998; Ratkiewicz et is such that in 2007 the shock should be several AU closer al., 1998; Pogorelov et al., 2004). Voyager 2 entered the ter- than average, suggesting that V2 might cross the shock in mination shock foreshock region much closer to the Sun 2007–2008. The models also indicate that the heliosheath (75 AU) than V1 (85 AU) (Stone et al., 2005). This large a is likely narrower at V2, so V2 may cross the heliopause at difference probably does not arise entirely from solar wind about the same time as V1. Voyager 1 and V2 observations dynamic pressure changes but also from asymmetries in the of the asymmetry of the termination shock and heliosheath boundary locations, with the boundaries closer in the direc- should make possible improved models and estimates of the tion of V2. distances to the heliopause. The of the heliosphere distortion depends on the orientation and strength of the interstellar magnetic field. 4. COSMIC RAYS Several suggestions for the magnetic field direction in the local interstellar cloud near the heliosphere have been pub- Cosmic rays are high-energy charged particles that bom- lished. The location of the sources of the heliospheric radio bard Earth from above the atmosphere. These particles have emissions (described below) lies along a line roughly par- important effects at Earth. Cloud cover is directly correlated allel to the galactic plane. Gurnett et al. (2006) suggest that with flux, so these particles may affect Earth’s the radio emission should occur on magnetic field lines tan- weather (Svensmark et al., 2007). Several thousand cosmic Richardson and Schwadron: The Limits of Our Solar System 449 rays pass through a person’s body every minute. These can ments are not neutral in the interstellar medium and there- cause biologic damage but also cause mutations that accel- fore cannot drift into the heliosphere. erate evolution. A new source of cosmic rays discussed be- Instruments such as the Solar Wind Ion Composition low is generated from the grains in the (Schwad- Spectrometer (SWICS) on Ulysses are able to detect pickup ron, 2002). These cosmic rays were likely more intense in ions directly. Ongoing measurements by cosmic-ray detec- the past and the termination shock was at times much closer tors on Voyager and Wind have detected additional ACR to the Sun than it is today. Both of these effects may have components (Reames, 1999; Mazur et al., 2000; Cummings caused periods when Earth was ensconced in much more et al., 2002). We now understand that, in addition to the severe radiation environments than at present. Several types traditional interstellar source, grains produce pickup ions of cosmic rays come from the outer heliosphere, as dis- throughout the heliosphere: Grains near the Sun produce cussed below. an “inner source” of pickup ions, and grains from the Kui- per belt provide an “outer” source of pickup ions and anom- 4.1. Galactic Cosmic Rays alous cosmic rays. Recent observations call into question not only the Galactic cosmic rays (GCRs) possess energies above sources of ACRs, but also the means by which they are ac- 100 MeV and are thought to be accelerated in the shock celerated. The prevailing theory, until V1 crossed the ter- waves associated with in our galaxy. The GCRs mination shock, was that pickup ions were energized at the are predominately fully stripped nuclei but lesser fluxes termination shock to the 10–100-MeV energies observed of GCR electrons are also observed. Galactic cosmic rays (Pesses et al., 1981). When V1 crossed the termination representing most elements in the periodic table have been shock, it did not see a peak in the ACR intensity as this the- detected. Since cosmic rays are charged, they are affected ory predicts (Stone et al., 2005). Instead, the ACR intensi- by magnetic fields in interplanetary space, the heliosphere, ties continued to increase in the heliosheath. One suggestion and near Earth. Although GCRs do reach Earth, many are is that the ACRs are accelerated in the flanks of the helio- deflected by the heliospheric magnetic field before they sheath and the source region is not observed by V1 near the reach the inner solar system. Fewer GCRs are observed at nose (McComas and Schwadron, 2006). Another sugges- Earth at solar maximum, when the solar wind magnetic field tion is that the acceleration region was affected by a series is stronger and more turbulent. Large merged interaction of MIRs before the V1 crossing that diminished the accel- regions (MIRs) temporarily reduce the flux of GCRs as eration process (McDonald et al., 2006; Florinski and Zank, these structures contain large and turbulent magnetic fields. 2006), in which case V2 may see a different ACR profile An empirical relation between GCRs and the magnitude of when it crosses the termination shock. B shows that the GCR flux increases when B is above nor- mal and increases when B is below normal; this relationship 4.3. The Interstellar and Inner Source Pickup Ions holds out to the termination shock (Burlaga et al., 2003a, 2005). Spacecraft in the outer solar system see increased The telltale signature of interstellar pickup ions, as seen GCR fluxes. These particles move so fast and are scattered for H+ in Fig. 4, is a cutoff in the distribution at ion speeds so thoroughly by the interstellar magnetic field that their twice that of the solar wind in the spacecraft reference flux is likely to be fairly uniform outside the heliosphere. frame. This cutoff is produced since interstellar ions are ini- tially nearly stationary in the spacecraft reference frame and 4.2. Anomalous Cosmic Rays subsequently change direction, but not energy, in this frame due to wave-particle interactions and gyration, causing them Early cosmic-ray observers discovered an unusual subset to be distributed over a range of speeds between zero and of cosmic rays that consisted of singly ionized ions (instead twice the solar wind speed in the spacecraft frame. of fully stripped nuclei) with energies of 1–50 MeV/nuc (see In comparison to the interstellar pickup ion distributions, Mewaldt et al., 1994). Most of the anomalous cosmic rays the C+ and O+ distributions in Fig. 4 are puzzling. The dis- (ACRs) are species that have high ionization thresholds, tributions do not cut off at twice the solar wind speed and such as He, N, O, Ne, and Ar. Until recently, ACRs were have a peak, contrary to the other ACR ions, which have thought to arise only from neutral in the interstellar flat distribution. These ions are singly charged, which pre- medium (Fisk et al., 1974) that drift freely into the helio- cludes them from having been emitted directly by the Sun; sphere through a process that has four essential steps: first, if this were the case, they would be much more highly the neutral particles from the local interstellar cloud enter the charged, e.g., C5+ and O6+ are each common solar wind heliosphere; second, the neutrals are converted into pickup heavy ions. The C+ and O+ distributions in Fig. 4 are consis- ions; third, the pickup ions are preaccelerated by shocks and tent with a pickup ion source close to the Sun. As the ions waves in the solar wind [see also section 5 and Schwadron travel out in the solar wind, they cool adiabatically in the et al. (1998)]; and finally, they are accelerated to their final solar wind reference frame, and instead of having a distri- energies at or beyond the termination shock (Pesses et al., bution that cuts off at twice the solar wind speed in the 1981). Easily ionized elements such as C, Si, and Fe are spacecraft reference frame, they have a peak near the solar expected to be strongly depleted in ACRs, since these ele- 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 geomet- ric cross-section typically 100 times larger than that inferred from 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 colli- sions of larger interplanetary grains (Wimmer-Schweingruber and Bochsler, 2003).

Fig. 4. Observed distributions from Ulysses/SWICS of C+ (solid 4.4. Pickup Ions from Cometary Tails triangles), O+ (open circles), and H+ (open triangles) vs. ion speed in the spacecraft frame normalized by the solar wind speed are also a source of pickup ions in the helio- (Schwadron et al., 2000). The observations are compared with sphere. Until recently, it was thought that the only way to + simulated distributions of solar wind H (dashed black curve), observe cometary pickup ions was to send spacecraft to + + interstellar H (upper gray dash-dot curve), inner source H (up- sample cometary matter directly. However, Ulysses has now per thick black line), inner source C+, O+ (lower thick black line), had three unplanned crossings of distant cometary tails, and interstellar O+ (lower gray dash-dot curve). 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 O+ data points is based on a transport model (Schwadron, three unplanned crossings of cometary tails by Ulysses sug- 1998; Schwadron et al., 2000) for ions picked up close to gest an entirely new way to observe comets through detec- the Sun, thereby proving the source is in the inner solar tion of distant cometary tails. system. Data like those presented in Fig. 4 suggested that another pickup ion source was present in addition to inter- 4.5. Outer Source Pickup Ions and stellar neutrals (Geiss et al., 1995; Gloeckler and Geiss, Anomalous Cosmic Rays 1998; 2000a; Schwadron et al., 2000). Since the source is peaked near the Sun, as is the distribution of interplanetary Recent observations from the Voyager and Wind space- grains, it was hypothesized that grains were associated with craft resolved ACR components comprised of easily ion- the source. ized elements [such as Si, C, Mg, S, and Fe (Reames, 1999; The initial, naive expectation was that the inner source Mazur et al., 2000; Cummings et al., 2002)]. An interstellar composition would resemble that of the grains, i.e., en- source for these “additional” ACRs, other than a possible hanced in C and O with strong depletions of noble elements interstellar contribution to C, is not possible (Cummings et such as Ne. However, the composition strongly resembled al., 2002). Thus, the source for these ACRs must reside that of the solar wind (Gloeckler et al., 2000a). This conun- within the heliosphere. A number of potential ACR sources drum was resolved by assuming a production mechanism are present within the heliosphere (Schwadron et al., 2002). whereby solar wind ions become embedded within grains The solar wind particles are easily ruled out as a source and subsequently reemitted as neutrals. Remarkably, the since they are highly ionized, whereas ACRs are predomi- existence of the inner source was hypothesized many years nantly singly charged. Discrete sources such as planets are prior to its discovery (Banks, 1971). also easily ruled out since their source rate is not sufficient The concept that the inner source is generated due to to generate the needed amounts of pickup ions. Another po- neutralization of solar wind requires that sputtered atoms tential source is comets, but the net cometary source rate is do not strongly contribute to the source. However, for grains sufficiently large only inside 1.5 AU, a location so close to larger than a micrometer, sputtering yields of the grains the Sun that the generated pickup ions would be strongly are much larger than the yields of neutralized solar wind cooled by the time they reach the termination shock, making (Wimmer-Schweingruber and Bochsler, 2003). This suggests acceleration to ACR energies extremely difficult. Moreover, that the grains that give rise to the inner source are ex- a cometary source would naturally be rich in C, which is tremely small (hundreds of angstroms). A small grain popu- not consistent with the compositional observations of easily lation generated through catastrophic collisions of larger ionized ACRs. interplanetary grains would also yield a very large filling The inner source (discussed in the preceding section) is factor (the filling factor is the net area of the sky filled in by 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 and due to wave-particle interactions; and finally, they are to the Sun. It has been suggested that the inner source may injected into an acceleration process at the termination shock be substantially preaccelerated in the inner heliosphere, to achieve ACR energies. The predicted abundances are all thereby overcoming the effects of adiabatic cooling. Con- within a factor of 2 of observed values, providing strong trary to this suggestion, inner source observations in slow validation of this scenario. See Schwadron et al. (2002) for solar wind show pronounced effects of adiabatic cooling a detailed presentation of these points. (Schwadron et al., 1999) and charge-state observations of energetic particles near 1 AU indicate no evidence for accel- 5. THE TERMINATION SHOCK FORESHOCK eration of the inner source (Mazur et al., 2002). The addi- tional population of ACRs requires a large source of pickup When V1 reached 84 AU in 2002 it observed increases in ions inside the heliosphere and outside 1 AU. energetic particles with energies from 40 keV to >50 MeV Schwadron et al. (2002) suggest that the source is ma- (Stone et al., 2005; Decker et al., 2005). Figure 6 shows that terial extracted from the Kuiper belt through a series of these ion fluxes increased in mid-2002 by a factor of 100, processes shown schematically by the lines in Fig. 5: First, disappeared in early 2003, came back at high energies in micrometer-sized grains are produced due to collisions of mid-2003 and at low energies in early 2004, and persisted objects within the Kuiper belt; grains spiral in toward the until the termination shock crossing in late 2004. These Sun due to the Poynting-Robertson effect; neutral atoms are intensity increases were not associated with solar events and produced by sputtering and are converted into pickup ions were not observed by V2, which trails V1 by 18.5 AU or when they become ionized; the pickup ions are transported about six years. Thus it seemed likely that these particles by the solar wind to the termination shock and, as they are were associated with the termination shock. But these par- convected, are preaccelerated due to interaction with shocks ticles 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 anisot- ropy 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 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) 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 direction led to suggestions that V1 had already crossed the the Sun (the actual field direction is tangential due to the termination shock (Krimigis et al., 2003) and that the par- winding of the Parker spiral with distance, so the particles ticles were coming outward from the termination shock. were moving tangentially but from the direction that is con- Subsequent work has shown that if the termination shock nected to the Sun). If the particle source were the termina- were blunt, the magnetic field lines first intersect the termi- tion shock and V1 was inside this shock, then the particles nation shock closer to the Sun at the distance of V1, so that seemed to be flowing the wrong way. This backward flow particles from the termination shock would flow to V1 in Richardson and Schwadron: The Limits of Our Solar System 453 the nominally outward direction (Jokipii et al., 2004). This data through mid-2007, these enhancements are not asso- hypothesis predicts that flows observed by V2, on the op- ciated with changes in the plasma or magnetic field. How- posite side of the nose of the termination shock, would be ever, when V1 was near the termination shock, in 2004, in the opposite direction (nominally toward the Sun). In late many of these enhancements were associated with cross- 2004, V2 observed these particles and they were in the op- ings of the heliospheric current sheet (HCS), but only from posite direction from those at V1, as predicted (Decker et negative to positive magnetic polarities (Richardson et al., al., 2006). Model predictions (Opher et al., 2006) show that 2006b). One explanation for this correlation is that particles field lines from the termination shock should intersect V1 from the termination shock current sheet drift (Burger and 2–3 AU upstream of the termination shock and V2 4–7 AU Potgieter, 1989) along the HCS to V1, and at southern to upstream of the shock. The amount of time each Voyager northern magnetic polarity HCS crossings the connection spends in the foreshock region depends on the motion of the to the termination shock is much closer. This mechanism termination shock. As V1 approached the termination shock would give peaks once per solar rotation. This hypothesis the solar wind dynamic pressure was increasing, pushing predicts that when V2, at southerly heliolatitudes, is close to the termination shock outward, so that V1 was surfing be- the termination shock, particle intensity peaks would be ob- hind the shock and remained in or near the foreshock for served only at positive to negative polarity crossings of the almost 9 AU. The solar wind pressure then decreased and HCS. Particle peaks are observed by both V1 and V2 away the shock moved quickly inward past V1. from the termination shock with periodicities of roughly a Figure 6 shows that the first termination shock particle solar rotation; the origin of these peaks is not understood. event ended at 2003.1. The end of this event coincided with the passage of a merged interaction region (MIR) by the 6. THE TERMINATION SHOCK CROSSING spacecraft that pushed the termination shock and thus the foreshock region outward. The first termination shock parti- The termination shock is where the solar wind plasma cle event at V2 also ended when a shock followed by a MIR makes its transition from supersonic to subsonic flow. The passed V2, driving the shock outward. The foreshock par- solar wind plasma is compressed and heated and the magnetic ticle fluxes observed have a lot of structure and many peaks. field strength increases. Zank (1999) reviews pre-encoun- Some of this structure is connected with the arrival of MIRs ter predictions of the nature and location of the termination at the location of V1 (Richardson et al., 2005b); since the shock. Since the average spiral magnetic field upstream of V1 plasma instrument does not work, the MIR arrival times the termination shock is along the T axis (where R is out- at V1 were estimated by propagating V2 data to the dis- ward from the Sun and T is in the solar equatorial plane and tance of V1. These MIRs can affect the particle fluxes in positive in the direction of the Sun’s rotation), the termina- several different ways; as in the above example, they can tion shock is expected to be a quasi-perpendicular shock. push the termination shock outward and sever or weaken However, energetic particles accelerated or reaccelerated the spacecraft connection to the shock. Since the MIRs push at the termination shock and in the turbulent heliosheath through the ambient solar wind, the magnetic field direc- can propagate into the upstream region. When the Voyager tion may change at the MIR or at the leading shock, thus spacecraft are near the termination shock and connected to changing the magnetic connection between the termination it by the magnetic field, measured particle intensities and shock and the spacecraft, which changes the observed parti- anisotropies may be comparable to those at the termination cle flux. Since MIRs are regions of enhanced magnetic field shock if pitch angle scattering is infrequent over the con- strength and turbulence, energetic particles diffuse through nection length and particles can traverse this length before these regions slowly and pile up ahead of the enhanced being convected shockward. Otherwise, intensities and an- magnetic field region, forming intensity peaks as the par- isotropies observed at Voyager will be reduced relative to ticles are “snowplowed” ahead of the MIRs (McDonald et those at the termination shock. al., 2000). Only one MIR has occurred since V2 entered the The termination shock was crossed in December 2004 on sheath; it was associated with an increase of lower-energy DOY 351, unfortunately during a tracking gap. The Voy- particles (28–540 keV) at the leading shock, probably due to ager spacecraft are tracked by the DSN 10–12 hours/day acceleration at the shock. The more-energetic 540–3500 keV due to competition with other spacecraft. The day before the particles had intensity peaks after the shock; these could shock crossing, the V1 Low Energy Charged Particle (LECP) result from snowplowing of particles ahead of the MIR. The instrument saw an intense beam of field-aligned ions and particle peak could also be behind the shock if the shock electrons at the same time as the Plasma Wave Spectrom- has weakened so that higher-energy particles are no longer eter (PWS) experiment saw electron plasma oscillations; accelerated; the particles already energized before the shock theory predicts that electron beams will drive these waves source turned off would convect with the solar wind while (Decker et al., 2005; Gurnett and Kurth, 2005). After the the shock would propagate ahead through the solar wind at shock crossing the lower-energy (tens of keV) ions jump in the fast mode speed (Decker et al., 2001). intensity by about an order of magnitude above those in the Both V1 and V2 observe periodic enhancements in the foreshock and the intensities become steady as shown in particle fluxes with periods of roughly a solar rotation. In Fig. 6. The termination shock particles shown in Fig. 7 are the first V1 foreshock encounter and in the V2 foreshock 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 (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 a strong (r = 4) and a weak (r = 2.4) shock (Cummings et al., 2002). celeration (Decker, 1988) theory for ions energized at a sistent with the magnetic field observations, which showed shock. The spectra of termination shock particles both in- the magnetic field strength increased by about a factor of side and outside the termination shock are best fit by two 3 at the shock and stayed high, convincing evidence that power laws with a break energy at 3.5 MeV. The constancy V1 was in the heliosheath (Burlaga et al., 2005). The stan- of the break energy suggests that the termination shock par- dard deviation of the magnetic field components increased ticles in the foreshock and heliosheath have the same source. across the termination shock, consistent with predictions for If the termination shock particles were accelerated by dif- the heliosheath region, and the distribution of the field mag- fusive shock acceleration, then the shock strength can be nitudes changed from log normal in the solar wind to determined from the slope of the termination shock particle Gaussian in the heliosheath. After V1 crossed the termina- power law and is 2.6 (2.4–3.0 with error bars) (Stone et al., tion shock into the heliosheath, the intensities of ions of all 2005). The shock strength is the compression ratio at the energies first increased, then fluctuated for 20–30 days, and shock (Ndownstream/Nupstream) and is 4 for a strong shock. An then became relatively steady, remaining essentially flat at alternative to diffusive shock acceleration has been sug- lower energies and increasing slowly at higher energies. At gested for the <20-MeV particles; a simple model propa- the termination shock crossing, the intensity of 40–53-keV gating the measured pickup ions from Ulysses outward and ions increased tenfold, reached a peak, then dropped and heating them at the termination shock can fit the observed thereafter was relatively smooth, varying typically by no spectra at these energies (Gloeckler et al., 2005). In this more than ≈20% about a mean of ≈300 for the next six model, up to 80% of the solar wind dynamic pressure goes months. The anisotropies went from beamlike in the fore- into heating the pickup ions. A relatively weak shock is con- shock to nearly isotropic in the heliosheath. Richardson and Schwadron: The Limits of Our Solar System 455

The intensity of the more-energetic ACR particles in- cosmic rays likely require substantial time to be acceler- creases at the shock but remains lower than the peak values ated; insufficient connection time near the nose could pre- in the foreshock. As V1 moved deeper into the heliosheath vent acceleration to ACR energies. The much longer con- the ACR intensities increased. These ions are isotropic, not nection times toward the distant flanks of the shock may flowing in field-aligned beams as in the foreshock. The ma- provide the necessary acceleration time to achieve the very jor surprise at the shock crossing was that the ACR inten- high 100–300-MeV ACR energies (McComas and Schwad- sities did not peak and the ACR spectra did not unroll into ron, 2006; Schwadron and McComas, 2007). If the termi- a power law distribution. For almost 30 years the paradigm nation shock has a strong nose-to-tail asymmetry, then the had been that ACRs were accelerated at the termination flanks would be connected by field lines that are well be- shock through diffusive shock acceleration. This theory pre- yond V1 near the nose, as illustrated in Fig. 8. dicted a peak in ACRs at the shock and that the acceleration Figure 9 documents the strong correlation between modu- process would result in a power law spectra at the termina- lated ACR and GCR helium. The data points show V1 obser- tion shock. Instead, the ACRs intensities at the termination vations (McDonald et al., 2006) from January 2002 to Febru- shock did not reach even the level in the foreshock and the ary 2006; the upper panel, energy range 150–380 MeV/ ACR energy spectra barely changed. Figure 7 compares the nuc, shows GCRs, and the lower panels, energy ranges 10– 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 ques- tion. 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 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 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. 9. The correlation between modulation of ACRs and GCRs. The panels show the relative fluxes of ACRs and GCRs observed Fig. 8. The configuration of the termination shock and magnetic by Voyager 1. The data (plus signs) were reduced by F. McDonald field in the heliosheath from Schwadron and McComas (2007). and B. Heikkila. All fluxes have been normalized by their maxi- The bluntness of the termination shock leads to ACR accelera- mum values and the fluctuations over time (averaged over 26 days) tion on the distant flanks (McComas and Schwadron, 2006; are plotted relative to one another. The respective energy ranges Schwadron and McComas, 2007). The red region shows the outer and species are indicated on the axes. The solid curves show theo- modulation boundary of ACRs accelerated on the distant flanks retical predictions based on a simple Parker-type diffusion-con- of the termination shock. vection 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 50 km/s in this time period, consistent with the termination normalized to the maximum value of the differential energy shock moving inward. flux in the respective energy ranges. Schwadron and Mc- An interesting feature in the magnetic field data are large Comas (2007) show how the strongly correlated modula- fluctuations in the field magnitude, a factor of 3, that are tion of ACRs and GCRs requires that the ACR accelerator purely compressional (the field direction does not change) be on field lines that connect far beyond V1 on either the (Burlaga et al., 2006). These features look qualitatively like distant reaches of termination shock flanks or far out in the the mirror mode waves observed in the magnetosheaths of inner heliosheath near the heliopause. and (Joy et al., 2006; Bavassano-Catteneo Another possibility for explaining ACR acceleration in et al., 1998); sometimes they are dips below the background the heliosheath is stochastic acceleration (Fisk, 2006). How- field value as are observed in low β regions of planetary ever, stochastic mechanisms encounter several problems magnetosheaths, and sometimes they are increase above the with explaining ACR acceleration at high energies: (1) the background field as observed in high β regions. The periods rates of stochastic acceleration are typically lower than the of these waves are roughly an hour. If they are mirror mode rate of diffusive acceleration at the termination shock, which waves, they result from an instability that occurs in high β makes it difficult to generate highly energetic ACRs in the plasmas when Tperp > Tpar, where Tperp and Tpar are the tem- one- to a few-year time frame before they become multiply peratures perpendicular and parallel to the magnetic field ionized (Jokipii, 1987; Mewaldt et al., 1996); and (2) the direction. These conditions often occur behind perpendicu- mechanism must lead to ACRs generated only far out in the lar shocks such as the termination shock. As Voyager gets heliosheath, or even toward the tail, whereas a distributed further from the shock, one expects that the plasma will, as process seems more likely to generate a very distributed a result of these waves, become more stable and the waves source. Although stochastic processes may not act quickly less frequent. enough to accelerate ACRs to high energies, they are criti- cal for generating the suprathermal seed population, which 7. THE HELIOSHEATH is naturally favored for diffusive acceleration at the termi- nation shock (Fisk, 2006). The spectra of the termination shock ions in the helio- McDonald et al. (2006) suggest that the lack of an ACR sheath are much less variable than in the foreshock. The peak was a time-dependence effect. A MIR passed V1 and steadiness of the spectrum suggests that the spectral shape is encountered the termination shock just before the termina- little affected by modulation in the heliosheath, as expected tion shock crossed V1; perhaps this reduced the accelera- (Jokipii and Giacalone, 2003). The energy spectrum in the tion efficiency and the ACR spectra recovered with time, heliosheath hardens (has a larger percentage of high-energy so the increase in intensity observed after the V1 termina- ions) with time; the intensities of the higher-energy ions tion shock crossing is a time, not radial distance, effect. This increase while those of the lower-energy ions are roughly scenario accounts for the larger ACR fluxes in the foreshock constant. The angular distributions of ions from 40 keV to than at the termination shock; the MIR reduces the ACR at least 30 MeV were highly anisotropic during the 2.5-year fluxes at the termination shock from the previous, larger approach of V1 to the termination shock. In the heliosheath intensities that were observed flowing into the foreshock the ion anisotropies are much smaller and more variable in from the termination shock. Over six months later, the ACR direction (Decker et al., 2005; Krimigis et al., 2005). fluxes were still recovering from this large temporal change. Figure 6 shows that 5–6 days after the termination shock This hypothesis will be tested when V2 crosses the termi- crossing, V1 began measuring a positive radial flow com- ≈ –1 nation shock. ponent fluctuating around VR +100 km s , which con- Another surprise was that after the termination shock tinued for 27–28 days. Near 2005.045, VR decreased rap- crossing V1 stayed in the same magnetic sector for 150 days idly (within about 5 days) and flow became inward from (Burlaga et al., 2005). Since the Sun rotates every 25.4 days 2005.063 to 2005.300. From about 2005.30 to 2005.4, VR –1 and the HCS is tilted, on average two HCS crossings should was 0 to +30 km s . VR increased again around 2005.45 occur every 25 days from southern magnetic polarities to and remained near +50 to +100 km s–1 until at least 2005.54. ≈ northern and back. These data were even more puzzling These excursions of VR over a period 90–140 days are since V1 was in the southern polarity sector, where V1 consistent with V1 sampling the variable flow downstream should spend less time since it is at northerly latitudes. Sec- of an inwardly moving termination shock. Convective flow tor crossings did resume (Burlaga et al., 2006); the expla- at V1 that fluctuates about zero implies that V1 would es- nation for the initial lack of crossings is that the termina- sentially be riding with and sampling the same parcel of tion shock and heliosheath were moving inward at a speed plasma for an extended period. Tangential flows are also ob- comparable to the outward heliosheath flow speed, so V1 served in the direction away from the heliospheric nose di- remained in the same heliosheath plasma as it moved out- rection, as expected for flow around the heliosphere (Decker ward (Jokipii, 2005). The solar wind pressure was decreas- et al., 2007). ing, which would make the termination shock and helio- The relationship between ACRs and termination shock sheath move inward. The radial speeds deduced from the particles is unknown. Both are deficient in C ions, indicating LECP data range from –30 to 100 km/s but average 30– that they are accelerated pickup ions (Krimigis et al., 2003). Richardson and Schwadron: The Limits of Our Solar System 457

However, the H/He ratio is ~10 for termination shock parti- The appearance of large variations in flow velocities associ- cles (Krimigis et al., 2003) and ~5 for ACRs (Cummings ated with disturbances from such relatively small-scale dy- et al., 2002), suggesting that He is more easily accelerated namical processes may indicate that V1 is near the helio- to ACR energies than H (Zank et al., 2001). pause. Reconnection at the heliopause-LISM interface could The flow streamlines in the heliosheath were calculated accelerate low-energy ions (and electrons) and create large by Parker (1963), assuming subsonic, incompressible, ir- departures from the normal heliosheath intensities and rotational flow. Using a similar flow model and a kinematic anisotropies. Intermittent reconnection along a slightly cor- approximation, Nerney et al. (1991) and Washimi and Ta- rugated heliopause surface might appear in the low-energy naka (1996) calculated the variation of the magnetic field ion (and possibly, electron) intensity-time profiles as grad- B in the heliosheath. Close to the termination shock, the ve- ual intensity increases with superposed, anisotropic inten- locity remains nearly radial and B remains nearly azimuthal. sity spikes as the spacecraft approached the heliopause, not The speed decreases as the plasma moves toward the helio- unlike structures encountered by V1 as it approached the pause; consequently, the magnetic field strength B in- termination shock. These comments are speculative, draw- creases. Cranfill (1971) and Axford (1972) showed that if ing largely upon analogies to planetary magnetopauses; the heliosheath is thick enough, B might become strong however, as with the termination shock, energetic particle enough to influence the flow. Nerney et al. (1993) estimated observations may enable us to remotely sense the helio- that B might influence the flow in a restricted region. The pause well before the Voyagers cross it. flow is deflected away from the radial direction as it ap- Instabilities on the heliopause have been the subject of proaches the nose of the heliopause, and B develops a radial theoretical conjecture (cf. Fahr, 1986; Florinski et al., component. At the heliopause, B must be parallel to the sur- 2005). In principle, the Rayleigh-Taylor instability could be face (unless significant reconnection occurs at the bound- important near the nose of the heliosphere where the dif- ary). The observations of B in the heliosheath show fea- ference between the flow densities are large. This instability tures not described by the simple kinematic models. While requires an effective “gravity” or destabilizing force. Liewer B does have a strong azimuthal component, a persistent me- et al. (1996) proposed that coupling between ions and neu- ridional component (northward) is observed that was not trals via charge exchange would produce such a destabiliz- predicted (Burlaga et al., 2005). ing force. Near the flanks, away from the nose, the - The magnetic field direction in the solar wind is often Helmholtz instability is more likely to be important where alternately toward and away from the Sun for several days. the shear velocity across the heliopause is large. Either type The existence of these “sectors” is related to extensions of of instability would produce motion of the heliopause su- fields from the polar regions of the Sun to the latitude of the perposed on the “breathing” expected to occur due to vary- observing spacecraft. Sectors have been observed by V1 out ing solar wind pressure. Estimates of the magnitude of the to 94 AU (Burlaga et al., 2003b, 2005). resulting excursions are tens of AU and have periods on the order of 100 years or more (Liewer et al., 1996; Wang and 8. THE HELIOPAUSE Belcher, 1998). For Voyager, the long periods make it un- likely that such wave motion can be detected, but the addi- As the Voyagers continue their trek through the un- tional radial motion from such instabilities, should they charted heliosheath, the next milepost in their journey will exist, increases the uncertainty in the distance to the helio- be the heliopause, which we think they will encounter be- pause. tween 2013 and 2021. The crossing of the heliopause will be a truly historic occasion; the first manmade vehicle to 9. THE HYDROGEN WALL leave the solar system. The heliopause is thought to be a tangential discontinuity with a large jump in the plasma Plate 11 shows that a stagnation point forms in the flow density and a rotation and change in the magnitude of the at the nose of the heliosphere. The plasma flow speed in magnetic field. Given the lack of an operational plasma in- this region slows, so the plasma is compressed and heated. strument on V1, the first crossing of the heliopause will Since the H+ and the H are coupled via charge exchange, probably be identified by a durable change in magnitude the H slows as well. Since the flux of H must be conserved, and/or direction of the magnetic field. If the model asym- slower H speeds translate to larger H densities (see Baranov metries discussed above are real, then the V1 and V2 cross- and Malama, 1993, 1995; Zank et al., 1996). Plate 11 shows ings could be close to each other in time. the region of denser H known as the hydrogen wall in front The flow velocity vectors are expected to become tan- of the heliosphere. This wall has been observed. Observa- gent to the nominal heliopause surface. However, variations tions of H Lyman-α emission from nearby stars in the nose of the heliopause surface and fluctuations in the flow ve- direction show that the H absorption is red-shifted relative locities could make it difficult to use these data to identify to the other lines by 2.2 km/s; its width is too broad to be a heliopause crossing. In analogy with planetary magneto- consistent with the T = 5400 ± 500 K temperature suggested pauses, the heliopause is probably a complex surface that by the width of the deuterium line (Linsky and Wood, 1996). varies locally in thickness and orientation and is the site of The profiles could only be fit assuming additional absorp- patchy reconnection and a variety of plasma instabilities. tion 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 100 km/s (Müller et al., 2006). Given possible values of the from the heliospheric hydrogen wall (which, since the H is interstellar pressure, in the past and future the heliopause moving sunward slower than the local interstellar cloud H, could be located anywhere between 12 and 400 AU and the produces absorption on the red side of the H line) and one termination shock from between 11 and 250 AU. For a very from the hydrogen wall of the nearby stars (which, since small heliosphere, cosmic-ray fluxes would dramatically in- the H is moving sunward faster than the local interstellar crease, affecting the climate and increasing rates of genetic cloud H, produces absorption on the blue side of the H line) mutation. When the heliosphere is in the , it (Linsky and Wood, 1996; Gayley et al., 1997; Izmodenov will be greatly expanded. The GCR fluxes at Earth will be et al., 1999, 2002; Wood et al., 1996, 2000). Many hydro- similar to present-day fluxes, but the fluxes of ACRS gen walls have been detected in front of other astrospheres, and neutral H will be greatly reduced. allowing comparison to the heliosphere and providing vali- dation to our general view of the heliosphere structure. 12. FUTURE OBSERVATIONS

10. THE HELIOSPHERIC RADIO EMISSIONS 12.1. Voyager

The Voyager spacecraft periodically detect radio emis- The Voyager spacecraft continue their journeys, V1 mov- sion from the boundaries of our solar system. This radio ing out about 3.5 AU/year and V2 about 3 AU/year. Voy- source is the most powerful in the solar system with an ager 2 is in the termination shock foreshock region; mod- emitted power of over 1013 W (Gurnett and Kurth, 1996). els suggest that this region only extends 4–7 AU from the Plate 12 shows a spectrogram of these emissions, plotting termination shock. In June 2007, V2 had already been in intensity as a function of time and frequency. These waves the foreshock region for almost 8 AU. The location of the have two frequency components, a constant frequency 2- termination shock and foreshock change with the solar wind kHz component and a component where the frequency rises pressure and thus the time spent in the foreshock can be from 2.4 to 3.5 kHz over a period of 180 days. The two longer than its width would suggest, as when the termina- largest events occurred in 1983–1984 and 1992–1994, a few tion shock was moving out just ahead of V1, but V2 will years after solar maxima. Large interplanetary shocks or probably cross the termination shock before the end of 2008. MIRs were hypothesized to trigger these emissions when This crossing will provide information on the asymmetry they entered the higher-density hydrogen wall beyond the of the heliosphere. It will also provide the first plasma data heliopause. But despite the occurrence of several large events from the heliosheath, providing direct information on the after the 2000 solar maximum, only very weak emission plasma flow, shock compression, and heating at the termi- was observed in 2002–2003, even though Voyager is pre- nation shock. Figure 10 shows the V1 and V2 trajectories sumably much closer to the radio source. The source re- in distance and heliolatitude. Model estimates of the loca- gion and trigger for these emissions is not well understood. tions of the termination shock and heliopause are shown shown by the lines. 11. THE HELIOSHEATH: THE LOCAL The V1 heliopause crossing is likely to occur between INTERSTELLAR MEDIUM 2013 and 2021 and the V2 heliopause crossing should occur in the same time frame. The spacecraft will have sufficient The local interstellar medium surrounding the helio- sphere 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 Fig. 10. Trajectories of the Voyager spacecraft in distance and within the past 40,000 years. Possible values of the densities heliolatitude. The range of predicted heliopause (for V1 and V2) encountered by the heliosphere in the past and future are and termination shock (for V2) distances are shown by the solid 0.005–15 cm–3 with Sun-cloud relative velocities of up to parts of the trajectory. Richardson and Schwadron: The Limits of Our Solar System 459

power to operate all instruments until 2016; after this time, lar medium. For example, Gloeckler et al. (2005) extrapo- power-sharing will extend the useful life of the spacecraft lated the V1 LECP proton spectra (40–4000 keV) upstream beyond 2020. Thus the Voyagers are likely to provide the and downstream of the termination shock to lower energies first in situ measurements of the local interstellar cloud. and demonstrated that 80% of the upstream ram pressure goes into heating the pickup protons. Thus V1 provides 12.2. the input proton spectrum for the global heliosheath (albeit only at the V1 location) that will be imaged in IBEX hy- New Horizons is a NASA mission that will fly by the drogen ENAs. And, returning the scientific favor, the very - system and then proceed to other Kuiper belt first IBEX half-year images will, at a glance, reveal asym- objects. Again, due purely to serendipity, Pluto is now lo- metries in the large-scale heliosheath configuration that bear cated toward the nose of the heliosphere. Although this on the puzzling directions of the magnetic field and ener- mission will not return data from the interstellar medium, getic particle streaming observed by V1 and V2. The IBEX it will measure the pickup ions directly, providing new in- all-sky images, in concert with local measurements from formation on the densities and composition of the interstel- Voyager, will provide a quantum leap forward in our un- lar neutrals entering the solar system. derstating of the outer heliosphere and its interstellar inter- action. 12.3. Interstellar Boundary Explorer 12.4. The The Interstellar Boundary Explorer (IBEX) is an excit- ing new NASA mission that will observe the large-scale Finally, in the era after Voyager and IBEX, we hope to structure of the heliosphere directly. It will be launched into send a new, optimally instrumented mission called the In- a highly eccentric Earth orbit in June 2008 [see the gen- terstellar Probe out beyond the termination shock, helio- eral description of the mission in McComas et al. (2004)]. pause, and bow shock and into the pristine interstellar me- All-sky energetic neutral (ENA) images of the global dium. Such a mission will require propulsion capable of structure of the termination shock and heliosheath will be accelerating a spacecraft to immense speeds if results are accumulated every 6 months. These ENAs are generated by to be gained in a reasonable mission lifetime. Recent studies charge-exchange collisions between the inflowing interstel- have examined both solar sails and nuclear-powered elec- lar neutral gas and the protons that are heated/accelerated tric propulsion. While the results of such a mission are still at the termination shock and that subsequently populate the decades off, we need to begin soon if we are to finally send entire heliosheath. Two IBEX single-pixel telescopes will a probe out into the material of the stars. these hydrogen ENAs over semiannually precess- ing great circles in the sky at low energies (10 eV to 2 keV) Acknowledgments. J.D.R. was supported under NASA con- and in an overlapping band of higher energies (300 keV to tract 959203 from the Jet Propulsion Laboratory to the Massachu- setts Institute of Technology. N.A.S. was supported by the NASA 6 keV). The intensities of the ENAs depend strongly on both Interstellar Boundary Explorer mission, which is part of the GSFC the shape of the proton spectrum and the plasma convec- Explorer Program, and by NASA’s Earth-- Radiation tion flow pattern in the heliosheath. 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