114 Encyclopedia of the Solar System
perpendicular to the field by a factor of∼1.1 on average at
1 AU. However, the temperature anisotropy for core elec-
trons varies systematically with density such that at very
low densities (<2cm−^3 ) the temperature ratio often ex-
ceeds 2.0, while at very high densities (>10 cm−^3 ) the tem-
perature ratio is often slightly less than 1.0. Such system-
atic variations of core electron temperature anisotropy with
plasma density reflect the marginally collisional nature of
the thermal electrons and their nearly adiabatic expansion
in the spiral magnetic field.
The suprathermal electrons consist of a beam of variable
width and intensity, known as the strahl, directed outward
from the Sun along the heliospheric magnetic field and a
more tenuous and roughly isotropic “halo” (Fig. 14b). The
angular width of the strahl results from a competition be-
tween focusing associated with conservation of magnetic
moment in the diverging heliospheric magnetic field and
defocusing associated with particle scattering. The strahl
carries the solar wind electron heat flux; variations in strahl
intensity largely reflect spatial variations in the corona from
which it arises. In addition, brief (hours) strahl intensifi-
cations often occur during solar electron bursts associated
with solar activity (see Section 7). The strahl serves as an
effective tracer of magnetic field topology in the interplan-
etary medium since its usual unidirectional nature arises
because field lines in the normal solar wind are “open” (see
Section 7.5) and are thus effectively connected to the so-
lar corona at only one end. In contrast, field lines thread-
ing ICMEs are often attached to the Sun at both ends (see
Section 7.4 and 7.5), and counterstreaming strahls are com-
monly observed there. Indeed, counterstreaming strahls are
one of the more reliable signatures of ICMEs (see Figs. 7b
and 9 and Table 2). Finally, the nearly isotropic electron
halo results primarily from the scattering out of the strahl
at distances beyond 1 AU and the subsequent reflection of
those backscattered electrons inside 1 AU by the stronger
magnetic fields that reside there.
11. Heavy Ion Content
Although the solar wind consists primarily of protons (hy-
drogen), electrons, and alpha particles (doubly ionized he-
lium), it also contains traces of ions of a number of heavier
elements. Table 3 provides estimates of the relative abun-
dances of some of the more common solar wind elements
summed over all ionization states. After hydrogen and he-
lium, the most abundant elements are carbon and oxygen.
The ionization states of all solar wind ions are “frozen in”
close to the Sun because the characteristic times for ioniza-
tion and recombination are long compared to the solar wind
expansion time. Commonly observed ionization states in-
clude He^2 +,C^5 +,C^6 +,O^6 +to O^8 +,Si^7 +to Si^10 +, and Fe^8 +
to Fe^14 +. Ionization state temperatures in the low-speed
wind are typically in the range 1.4 to 1. 6 × 106 K, whereas
ionization state temperatures in the high-speed wind are
TABLE 3 Average Elemental Abundances in the
Solar Wind
Element Abundance Relative to OxygenH 1900 ± 400
He 75 ± 20
C 0.67±0.10
N 0.15±0.06
O 1.00
Ne 0.17±0.02
Mg 0.15±0.02
Si 0.19±0.04
Ar 0.0040±0.0010
Fe 019 +0.10,−0.07typically in the range 1.0 to 1. 2 × 106 K. Unusual ionization
states such as Fe+^16 and He+^1 , which are not common in
the normal solar wind, are often abundant within ICMEs,
reflecting the unusual coronal origins of those events.
The relative abundance values in Table 3 are long-
term averages; however, abundances vary considerably
with time. Such variations have been extensively studied
for the He^2 +/H+ratio, A(He), but are less well established
for heavier elements. The most probable A(He) value is
∼0.045, but the A(He) ranges from less than 0.01 to val-
ues of 0.35 on occasion. The average A(He) is about half
that commonly attributed to the solar interior, for reasons
presently unknown. Much of the variation in A(He) and in
the abundance of heavier elements is related to the large-
scale structure of the wind. For example, Fe/O and Mg/O
ratios are systematically lower in high-speed streams than
in low-speed flows. A(He) tends to be relatively constant
at∼0.045 within quasi-stationary, high-speed streams but
tends to be highly variable within low-speed flows. Particu-
larly low (<0.02) abundance values are commonly observed
at the heliospheric current sheet. A(He) values greater than
about 0.10 are relatively rare and account for less than 1%
of all the measurements. At 1 AU, enhancements in A(He)
above 0.10 occur almost exclusively within ICME plasma.
The physical causes of these variations are uncertain for the
most part, although thermal diffusion, gravitational settling,
and Coulomb friction in the chromosphere and corona all
probably play roles.12. Energetic Particles
A proton moving with a speed of 440 km/s has an energy
of∼1 keV. Thus, by most measures, solar wind ions are
low-energy particles. The heliosphere is, nevertheless, filled
with a number of energetic ion populations of varying in-
tensities with energies ranging upwards from∼1to∼ 108
keV/nucleon. These populations include galactic cosmic
rays, anomalous cosmic rays (see discussion that follows),