112 Encyclopedia of the Solar System
parallel to the heliopause. The shape of the heliosphere
is asymmetric because of its motion relative to the inter-
stellar gas; it is compressed in the direction of that motion
and is greatly elongated in the opposite direction. Observa-
tions in the outer heliosphere suggest that the termination
shock is constantly in motion relative to the Sun owing to
an ever-changing solar wind momentum flux; it may never
truly achieve an equilibrium position. The size and shape
of the heliosphere depend on the momentum flux carried
by the solar wind, the pressure of the interstellar plasma,
and the motion of the heliosphere relative to the interstel-
lar medium.Voyager 1recently verified the existence of the
termination shock, having crossed it in December 2004 at
a heliocentric distance of 94 AU roughly in the direction of
the heliosphere’s motion relative to the interstellar medium.
It is currently believed that the heliopause lies at a helio-
centric distance of 115–150 AU and should be encountered
by Voyager within the next decade.
10. Kinetic Properties of the Plasma
10.1 The Solar Wind as a Marginally
Collisional Plasma
On a large scale, the solar wind behaves like a compressible
fluid and is capable of supporting relatively thin structures
such as shocks. It is perhaps not obvious why the solar wind
should exhibit this fluid-like behavior since the wind is a
dilute plasma in which collisions are relatively rare. For ex-
ample, using values given in Table 1, we find that the time
between collisions for a typical solar wind proton at 1 AU is
several days. (These collisions do not result from direct par-
ticle impacts such as colliding billiard balls, but rather from
the long distanceCoulomb interactionscharacteristic of
charged particles.) The time between collisions is thus com-
parable to the time for the solar wind to expand from the
vicinity of the Sun to 1 AU; this is the basis for statements
that the solar wind is a marginally collisional plasma.
There are several reasons why the solar wind behaves
like a fluid even in the absence of particle collisions to ef-
fect fluid-like behavior. First, when the temperature is low
and the density is high, collisions are more frequent than
noted previously. Second, the presence of the heliospheric
magnetic field causes charged particles to gyrate about the
field, and they thus do not travel in straight lines between
collisions. For typical conditions at 1 AU, solar wind elec-
trons and protons havegyro radiiof∼1.4 and∼60 km,
respectively, which are small compared to the scale size of
most structures in the solar wind. Third, the solar wind is
subject to a variety of instabilities that are triggered when-
ever particle distribution functions depart significantly from
thermal distributions (see Section 10.2). These instabilities
produce collective interactions that mimic the effects of par-
ticle collisions. Finally, because the magnetic field is frozen
into the solar wind flow, parcels of plasma originating from
different positions on the Sun cannot interpenetrate one
another.10.2 Kinetic Aspects of Solar Wind Ions
Collisional gases can usually be described by a single
isotropic (i.e., the same in all directions) temperature
(T) with the distribution of particle speeds (v) obeying
f(v)∼exp[−m(v−V 0 )^22 kT], where f is the number of
particles per unit volume of velocity space,kis Boltzman’s
constant (1.38× 10 −^16 erg/deg),mis the particle mass, and
V 0 is the bulk speed of the gas. In contrast, proton distribu-
tion functions in the solar wind are usually anisotropic be-
cause of the paucity of collisions and because the magnetic
field provides a preferred direction in space. At 1 AU, the
proton temperature parallel to the field is generally greater
than the temperature perpendicular to the field, on aver-
age by a factor of∼1.4. Moreover, solar wind proton and
alpha particle distributions often exhibit significant non-
Maxwellian features such as the double-peaked distribu-
tions illustrated in Fig. 13a. The secondary proton and alpha
particle peaks are associated with beams streaming relative
to the main solar wind component along the heliospheric
magnetic field. The relative streaming speed of such beams
is usually comparable to or less than the local Alfv ́en speed,
suggesting that the streaming is limited by a kinetic insta-
bility. Closer to the Sun, where the Alfv ́en speed is higher,
relative streaming speeds between the beams and the main
components can be as large as several hundred km/s. Sec-
ondary proton beams are common in the solar wind in both
low- and high-speed flows and may play a fundamental role
in the overall acceleration and heating of the wind; how-
ever, their origin in solar and/or heliospheric processes is
presently uncertain. Figure 13b illustrates that solar wind
ion distributions in the low-speed wind also commonly have
extended nonthermal tails of uncertain origin. Particles in
these extended tails are easily accelerated to much higher
energy when they encounter shocks (see Section 7).10.3 Kinetic Aspects of Solar Wind Electrons
Electron distributions in the solar wind consist of a rel-
atively cold and dense thermal “core” population that is
electrically bound to the solar wind ion population and a
much hotter and freer-running suprathermal population
that becomes collisionless close to the Sun. At 1 AU, the
breakpoint between these populations typically occurs at
∼70 eV (Fig. 14a). This breakpoint moves steadily to lower
energies with increasing heliocentric distance as the core
population cools. Typically the core contains about 95% of
the electrons, and at 1 AU has a temperature of∼1.3× 105.
The core electrons typically are mildly anisotropic, with the
temperature parallel to the field exceeding the temperature