The structure of the heliosphere

2.3 The structure of the heliosphere

Wind/ISM interactions provide the means by which solar-like stellar winds can be detected (see Section 4). Our understanding of these interactions relies heavily on a long history of efforts to model the solar wind/ISM interaction. This heliospheric modeling is summarized briefly here, but for more comprehensive reviews see Holzer (1989 ); Baranov (1990 ); Suess (1990 ) and Zank (1999 ).

Modeling the large scale structure of the heliosphere began not long after the solar wind’s discovery (Parker, 1961 , 1963 ). The basic structure of the heliosphere, which is shown schematically in Figure 3 , is dominated by three prominent boundaries: the termination shock (TS), heliopause (HP), and bow shock (BS). The solar wind is highly supersonic, and the oval-shaped termination shock is where the radial wind is shocked to subsonic speeds. The $26 km s− 1$ laminar ISM flow is also generally believed to be supersonic, although in principle it could be subsonic if the poorly known ISM magnetic field is strong enough (Zank et al., 1996 ). Nevertheless, most heliospheric models assume the existence of a bow shock, where the ISM flow is shocked to subsonic speeds (see Figure 3 ). In between the TS and BS is the heliopause, which is a contact surface separating the plasma flows of the solar and interstellar winds.

Figure 3:

Schematic picture of the heliospheric interface from Izmodenov et al. (2002

), which can be divided into the 4 regions shown in the figure, with significantly different plasma properties. Region 1: supersonic solar wind; Region 2: subsonic solar wind; Region 3: disturbed interstellar gas and plasma; and Region 4: undisturbed interstellar medium.

The heliospheric structure shown in Figure 3 is inferred almost entirely from hydrodynamic models. However, in 2004, Voyager 1 crossed the TS at a distance of 94 AU from the Sun in roughly the upwind direction of the ISM flow (Stone et al., 2005 ). Precursors of this crossing were seen years prior to the accepted crossing date (Krimigis et al., 2003 ; Burlaga et al., 2003 ; McDonald et al., 2003 ). At the time of this writing, Voyager 1’s sister satellite Voyager 2 has begun to see these precursors but has not yet officially crossed the TS (Opher et al., 2006 ). The 94 AU TS distance measured by Voyager 1 is consistent with model predictions (Izmodenov et al., 2003 ). As for the HP and BS, recent models generally predict upwind distances of $∼ 140AU$ and $∼ 240AU$ , respectively. The Voyager satellites may not survive long enough to get out this far. Update

The plasma component of the LISM is diverted around the heliopause due to the strong plasma interactions, but neutrals in the LISM can penetrate into the solar system through the HP and TS. These neutrals were first detected through Ly $α$ backscatter emission (Bertaux and Blamont, 1971 ). However, even after this discovery most hydrodynamic models of the heliosphere continued to ignore the neutrals since the collisional interactions involving neutrals are much weaker than those involving charged particles. Essentially, the assumption was made that the neutrals would pass through the heliosphere unimpeded, having little or no effect on the heliospheric structure.

It was recognized in the 1970s that the LISM neutrals could in fact play an important role in the solar wind/ISM collision through charge exchange interactions (Holzer, 1972 ; Wallis, 1975 ). However, trying to model this is very difficult, because the charge exchange sends the neutral H wildly out of thermal and ionization equilibrium. This means that simple fluid approximations break down and one has to resort to complex multi-fluid codes or ideally fully kinetic codes. It was not until much later that the first codes that treat the plasma and neutrals in a self-consistent manner were first developed (Baranov and Malama, 1993 , 1995 ; Baranov and Zaitsev, 1995 ; Zank et al., 1996 ; Izmodenov et al., 1999a ; Müller et al., 2000 ; Izmodenov et al., 2001 ). These models demonstrate that the heliospheric structure is indeed influenced significantly by the neutrals in many different ways.

Figure 4:

(a) Proton temperature, (b) proton density, (c) neutral hydrogen temperature, and (d) neutral hydrogen density distributions for a heliospheric model from Wood et al. (2000b

). The positions of the termination shock (TS), heliopause (HP), and bow shock (BS) are indicated in (a), and streamlines indicating the plasma flow direction are shown in (b). The distance scale is in AU.

For purposes of this article, the importance of the development of heliospheric codes that treat neutrals properly is that it is only because of the existence of neutral hydrogen in the LISM that heliospheric and astrospheric Ly $α$ absorption is detectable, and it is only because of the existence of the self-consistent codes developed to model neutrals in the heliosphere that we can model the astrospheric absorption and extract stellar mass loss rates from the data. Figure 4 shows a heliospheric model that uses a hybrid kinetic code, where the protons are modeled as a fluid but a kinetic code is used for the neutrals (Lipatov et al., 1998 ; Müller et al., 2000 ; Wood et al., 2000b ). The strong plasma interactions heat and compress LISM protons in between the HP and BS, and thanks to charge exchange processes these high temperatures and densities are transmitted to the neutral H. As a consequence, the heliosphere and astrospheres are permeated by a population of hot hydrogen, which produces a substantial amount of Ly $α$ absorption in HST observations of nearby stars. Most of this absorption comes from the “hydrogen wall” region in between the HP and BS, where densities of the hot H I are particularly high (see Figure 4 d). The heliospheric and astrospheric Ly $α$ diagnostic is described in detail in Section 4.