The Solar Wind
J. T. Gosling
LASP, University of Colorado
June 11, 2009
A Brief Overview
The solar wind is a plasma, i.e., an ionized gas, that fills the solar system.
It results from the supersonic expansion of the solar corona.
The solar wind consists primarily of electrons and protons with a smattering
of alpha particles and other ionic species at low abundance levels.
At 1 AU (Earth) average proton densities, flow speeds and temperatures are
~8.7 cm-3, 468 km/s, and 1.2 x 105 K, respectively.
Embedded within the solar wind is a magnetic field having an average
strength ~6.2 nanotesla at 1 AU.
The solar wind plays an essential role in shaping and stimulating planetary
magnetospheres and ionic comet tails. It is the prime source of space
weather.
Early Indications of the Solar Wind
Carrington’s 1859 observation of white light solar flare, followed 17 hours
later by a large geomagnetic storm - suggested possible cause and effect.
Lindemann (early 1900s) suggested large geomagnetic storms resulted
from interaction between Earth’s magnetic field and plasma clouds ejected
from Sun during flares.
Observations of recurrent (at 27-day rotation period of Sun) geomagnetic
storms led to hypothesis of M (for magnetic) regions on Sun that produced
long-lived streams of charged particles in interplanetary space.
There almost always is at least a low level of geomagnetic activity. This
suggested that plasma from the Sun is always present near Earth.
Observations by S. Forbush in 1930s and 1940s of modulations of cosmic
rays in association with geomagnetic storms and in association with 11year solar activity cycle suggested that the modulations were caused by
magnetic fields embedded in plasma clouds from the Sun.
Biermann concluded in early 1950s that a continuous outflow of particles
from the Sun filling interplanetary space was required to explain the antisunward orientation of ionic comet tails.
Parker’s Solar Wind Model
In 1958, motivated by diverse indirect observations, E. N. Parker developed the first
fluid model of a continuously expanding solar corona driven by the large pressure
difference between the solar corona and the interstellar plasma. His model produced
low flow speeds close to the Sun, supersonic flow speeds far from the Sun and
vanishingly low pressures at large heliocentric distances. In view of the fluid
character of the model, he called this continuous supersonic expansion the “solar
wind”.
Parker’s Model of the Heliospheric Magnetic Field
The electrical conductivity of the solar wind plasma is so high that the solar
magnetic field is “frozen into” the solar wind flow as it expands outward from
the Sun. Because the Sun rotates with a period of 27 days as observed from
Earth, magnetic field lines in the Sun’s equatorial plane are bent into spirals
whose inclination to the radial direction depend on heliocentric distance and
the speed of the wind. At 1 AU the average field is inclined ~45˚ to the radial
direction in the equatorial plane.
Axes are heliocentric
distance in units of AU.
First Direct Measurements of the Solar Wind
Provided Confirmation of Parker’s Basic Model
Measurements made by an electrostatic
analyzer and a magnetometer onboard
Mariner II during its epic 3-month
journey to Venus in 1962 provided firm
confirmation of a continuous solar wind
flow and spiral heliospheric magnetic
field that agree with Parker’s model, on
average.
Mariner II also showed that the solar
wind was highly variable, being
structured into alternating streams of
high and low-speed flows that lasted for
several days each. The observed
magnetic field was also highly variable
in both strength and orientation.
The Variable Solar Wind at 1 AU
n is proton density, Vsw is solar wind speed, B is magnetic field strength, A(He) is
He++/H+ ratio, Tp is proton temperature, Te is electron temperature, T is alpha
particle temperature, Cs is sound speed, CA is Alfven speed.
The Sun yearly loses ~6.8 x 1019 g to the solar wind, a very small fraction
of the total solar mass of ~2 x 1033 g.
Coronal and Solar Wind Stream Structure
The corona is highly non-uniform, being
structured by the interplay between the complex
solar magnetic field and the outflow of the solar
wind. It is thus not surprising that the solar wind
also is highly structured.
The recurrent high-speed streams originate in
coronal holes, which are large nearly unipolar
magnetic regions of low plasma density.
Low-speed flows tend to originate in coronal
streamers which straddle regions of field polarity
reversals in the solar atmosphere.
Each high-speed stream is asymmetric (rapid
rise, slower fall) and unipolar (Br positive or
negative) throughout. Reversals in Br occur in
the low-speed wind.The field strength and
plasma density (not shown) peak on the leading
edges of the streams, and the flow there is
deflected first westward (positive flow azimuth)
and then eastward.
The Heliospheric Current Sheet and the Solar Dipole
The Sun’s large-scale magnetic field well above the photosphere is usually
well approximated by that of a dipole.
The dipole generally is tilted relative to the rotation axis of the Sun, the tilt
changing with the advance of the 11-year solar activity cycle.
The heliospheric current sheet separates solar wind regions of opposite
magnetic polarities and wraps entirely around the Sun. It is the extension of
the solar magnetic equator into the heliosphere.
When the Sun’s magnetic dipole is tilted relative to the rotation axis, the
heliospheric current sheet is warped and resembles a ballerina’s twirling
skirt. In this simple picture, each ridge in the skirt corresponds to a different
solar rotation; the ridges are separated radially from one another by about
4.7 AU.
Solar Latitude Effects
Ulysses is in a solar orbit that took it to heliographic latitudes of +/- 80˚ in its
5.5-year journey about the Sun.
During the decline of solar activity and near solar minimum stream structure
is confined to a relatively narrow latitude band centered on the solar equator
because: 1) solar wind properties change rapidly as a function of distance
from the heliospheric current sheet; and 2) the current sheet is usually found
within ~+/- 30˚ of the solar equatorial plane at this phase of the solar cycle.
Near solar activity maximum
solar wind variability extends up
to the highest solar latitudes
sampled by Ulysses, as does
coronal structure.
Evolution of Stream Structure with Heliocentric Distance
Spatial variability of the solar wind outflow and
solar rotation produce radial variations in
speed.
R
F
Faster wind overtakes slow wind ahead while
outrunning slow wind behind. As a result, the
leading edges of high-speed streams steepen
with increasing heliocentric distance.
Plasma is compressed on the leading edge of a
stream and rarefied on the trailing edge.
The build up of pressure on the leading edge of
a stream produces forces that accelerate the
low-speed wind ahead and decelerate the highspeed wind within the stream
When the difference in speed between the crest of a stream and the
trough ahead is greater than about twice the sound speed, ordinary
pressure signals do not propagate fast enough to keep the stream from
“toppling over” and a forward-reverse collisionless shock pair forms on the
opposite sides of the high-pressure region to prevent that.
Evolution of Stream Structure with Heliocentric Distance
(continued)
R
F
Although the shocks propagate in opposite
directions relative to the solar wind, both
are carried away from the Sun by the
highly supersonic flow of the wind.
The major accelerations and decelerations
of the wind then occur at the shocks and
the stream profile becomes a damped,
double sawtooth.
Because the sound speed decreases with
increasing heliocentric distance, virtually all
high-speed streams eventually have shock
pairs on their leading edges.
The dominant structure in the solar
equatorial plane in the outer heliosphere is
the expanding compression regions where
most of the plasma and magnetic field are
concentrated.
Stream Evolution in Two and Three Dimensions
When the coronal expansion is spatially
variable but time-stationary, a steady flow
pattern such as shown here develops in the
equatorial plane.
The pattern co-rotates with the Sun and the
compression region is known as a
corotating interaction region, CIR.
Only the pattern rotates; each parcel of
solar wind plasma moves nearly radially
outward.
The compression region is nearly aligned with the magnetic field line spirals
and the pressure gradients thus have both radial and transverse components.
Thus the slow wind gets deflected to the west (left) and the fast wind gets
deflected to the east (right), as illustrated in slide #8.
*It is both observed and predicted that the CIRs typically have substantial
opposed meridional tilts in the opposite solar hemispheres .
Damped High-Speed Streams in the Outer Heliosphere
Voyager 2 data obtained at ~18 AU
Stream amplitudes are strongly damped in the outer heliosphere because
of the interactions between high and low-speed streams.
Solar Wind Electrons
Measurements of electron energy distributions in the solar wind reveal the
presence of both thermal and suprathermal populations.
The suprathermal population is nearly collisionless, carries the solar wind heat
flux, and includes both a field-aligned “strahl” (or beam) and a roughly
isotropic “halo”.
The suprathermal electrons behave as extremely fast test particles and serve
as very effective tracers of magnetic field line topology in the solar wind.
Coronal Mass Ejections and Transient Solar Wind
Disturbances
The most dramatic temporal evolution in the corona occurs in coronal
mass ejection events, CMEs, which, in turn, produce the largest transient
disturbances in the solar wind. The shock ahead of a fast CME is broader
than the CME that drives it. The ambient magnetic field drapes about the
CME.
Counterstreaming Suprathermal Electrons as
Tracers of Closed Magnetic Field Lines in CMEs
In the normal solar wind field lines are open to the outer boundary of the
heliosphere and a single field-aligned, anti-sunward-directed strahl is
observed.
CMEs originate in closed field regions in the corona and field lines within
CMEs are at least initially connected to the Sun at both ends.
Counterstreaming strahls are commonly observed on closed field lines
and help identify CMEs in the solar wind (ICMEs).
The Magnetic Field Topology of CMEs and the
Problem of Magnetic Flux Balance
3D Magnetic Reconnection Within the
Magnetic Legs of a CME
Possible mixture of Resulting
Field Topologies
Every CME carries new magnetic flux into the heliosphere. Magnetic
reconnection in the footpoints serves to open up the closed field loops
associated with a CME, produces helical field lines within it, and helps to
maintain a roughly constant magnetic flux in the heliosphere.
A Simple 1D Fluid Simulation of a Solar Wind
Disturbance Driven by a Fast CME
F
R
The simulation was initiated by raising
the flow speed from 275 to 980 km/s for 6
hours at the inner boundary.
A region of high pressure develops on
the leading edge of the disturbance as
the CME overtakes the slower wind
ahead. The high pressure compression
region is bounded by a forward shock, F,
on its leading edge and a reverse shock,
R, on its trailing edge. Momentum is
transferred from the fast CME to the
slower wind ahead via the F shock and
the CME slows with increasing
heliocentric distance. The reverse shock
is only rarely actually observed.
A Solar Wind Disturbance Produced by a
Moderately Fast Coronal Mass Ejection, CME
270 eV
electrons
Shock
CME
From top to bottom the quantities are
the pitch angle distribution (relative to
B) of suprathermal electrons, proton
density, proton temperature, speed,
A(He), field strength, and field angles.
The CME had a higher speed than that
of the ambient wind ahead of the shock
and in this case is distinguished by
counterstreaming suprathermal
electrons, anomalously low proton
temperatures, an elevated helium
abundance, and a relatively strong,
smoothly rotating, magnetic field.
CMEs in the Solar Wind
(ratio of gas pressure to magnetic field pressure)
Commonly Observed Ionization States in the Solar Wind
He2+
C5+, C6+
O6+ to O8+
Ionization states are “frozen in” close to the Sun
because the characteristic times for ionization
and recombination are long compared to the
solar wind expansion time.
Si7+ to Si10+
Fe8+ to Fe14+
Ionization state temperatures reflect electron temperatures in the solar corona
where the ionization states freeze in and are typically 1.4 - 1.6 x 106 ˚K in the
low-speed wind and 1.0 - 1.2 x 106 ˚K in the high-speed wind. Note that this
speed/temperature relationship is opposite to that predicted by Parker.
Unusual ionization states such as He+1 and Fe+16 are relatively common in
ICMEs, reflecting the unusual coronal origins of those events.
Turbulence and Alfven Waves in the Solar Wind
The solar wind is filled with fluctuations that have their largest amplitudes
in the high-speed wind. Many of these fluctuations are Alfvenic in nature
(coupled changes in flow velocity and magnetic field vectors). The Alfvenic
fluctuations are probably remnants of waves and turbulence that heat the
corona and accelerate the solar wind. Fluctuation amplitudes decrease
with increasing heliocentric distance; their dissipation heats the wind far
from the Sun.
Local Magnetic Reconnection in the Solar Wind
The solar wind contains numerous current sheets where the magnetic
field orientation changes abruptly. When magnetic reconnection occurs at
these current sheets the magnetic topology changes and oppositely
directed jets of plasma are produced. Observations of this jetting plasma
and the related magnetic field structure in the solar wind provide important
information about the reconnection process and its after-effects in
collisionless plasmas.
Interaction of the Solar Wind with the Interstellar Medium
The solar wind carves a cavity in the local interstellar plasma since the two
plasmas cannot readily interpenetrate one another. The size and shape of the
cavity depend on the momentum flux carried by the solar wind, the pressure of
the interstellar plasma, and the motion of the Sun relative to the interstellar
medium. In the last few years both Voyager 1 and Voyager 2 crossed the
termination shock, where the solar wind was substantially slowed, deflected,
and heated.
Energetic Particles in the Solar Wind
The heliosphere is filled with a variety of energetic ion populations of
varying intensities with energies ranging from ~1 to 108 keV/nucleon. Most
of these populations are the result of particle acceleration at shocks.
The Solar Wind as a Natural Plasma Laboratory
One of the first great triumphs of the space age was
experimental proof of the existence of a solar wind that fills the
solar system.
The solar wind serves as a magnificent natural laboratory for
studying and obtaining understanding of processes and
phenomena that occur in a variety of other space plasma and
astrophysical contexts. These include at least the following:
Kinetic, fluid and MHD aspects of plasmas
Plasma heating and acceleration
Collisionless shock physics
Energetic particle production and transport
Magnetic reconnection
Evolution and dissipation of waves and turbulence
Some Recent Hot Topics in Solar Wind Research
Magnetic reconnection in the solar wind
Physics of the termination shock and heliosheath
Origin of the low-speed wind - role of interchange reconnection
New ideas about nature/origin of heliospheric magnetic field
Magnetic field topology and flux balance in the heliosphere
Ionic composition variations
Turbulence dissipation and plasma heating
Energetic particles: production, sources and propagation
Heating and acceleration of the solar wind
CME origins and evolution in the heliosphere
The pickup of newly borne ions and their sources