The Solar Wind J. T. Gosling LASP, University of Colorado Boulder, CO

June 16, 2011

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 10 5 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.

Some 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 always is at least a low level of geomagnetic activity. This suggested to Birkeland that flowing 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 11 year 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 anti sunward 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 of the corona 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, V sw A(He) is He ++ /H + ratio, T is solar wind speed, B is magnetic field strength, p is proton temperature, T e is electron temperature, T Alfven speed.

 is alpha particle temperature, C s is sound speed, C A is The Sun yearly loses ~6.8 x 10 19 g to the solar wind, a very small fraction of the total solar mass of ~2 x 10 33 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 expansion of the solar wind. It is thus not surprising that a spacecraft observes a varying solar wind as the Sun rotates (once every 27 days as observed at Earth).

The Heliospheric Current Sheet and the Solar Dipole

The Sun’s large-scale magnetic field well above the photosphere is usually reasonably well approximated as 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, HCS, 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 HCS 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 and Solar Cycle Effects

During the decline of solar activity and near solar minimum solar wind variability is confined to a relatively narrow latitude band centered on the solar equator because: 1) solar wind speed increases rapidly with distance from the heliospheric current sheet (HCS); and 2) the HCS 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.

Characteristics of Solar Wind Stream Structure

1-hr averaged data Each high-speed stream is asymmetric (rapid rise, slower fall) and unipolar throughout. Reversals in field polarity occur at the HCS in the low-speed wind. The field strength and plasma density and temperature peak on the leading edges of the streams, and the flow there is deflected first westward (positive flow azimuth) and then eastward.

Evolution of Stream Structure with Heliocentric Distance

R F Spatial variability of the solar wind outflow and solar rotation produce radial variations in speed.

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.

Pressure gradient forces on the leading edge of a stream accelerate the low-speed wind ahead and decelerate the high-speed 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.

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

.

Stream Evolution in Two 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 originally slow wind in a CIR gets deflected to the west (left) and the originally fast wind in a CIR gets deflected to the east (right).

CIRs also have an interesting 3D structure – not discussed here.

Solar Wind Electrons

Measurements of electron velocity distributions in the solar wind reveal the presence of both thermal and suprathermal populations.

The suprathermal population is nearly collisionless and includes both a field aligned “strahl” (beam) resulting from focusing in the diverging solar wind magnetic field and a roughly isotropic “halo” resulting from scattering.

The suprathermal electrons are extremely fast 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 changes in the corona occur in coronal mass ejections, 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).

A Simple 1D Fluid Simulation of a Solar Wind Disturbance Driven by a Fast CME

R F 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 is bounded by a forward shock on its leading edge and a reverse shock on its trailing edge. Momentum is transferred from the fast CME to the slower wind ahead and behind by pressure gradients, and the CME slows and expands with increasing heliocentric distance.

The reverse shock is generally absent in 3D simulations.

3D Simulation of CME-Driven Disturbances in a Simply Structured, Tilted Dipole Solar Wind CME-driven disturbances can become highly distorted as they interact with ambient structure in the solar wind.

CMEs in the Solar Wind

(ratio of gas pressure to magnetic field pressure)

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 close to the Sun 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.

Variation of Solar Rotation-Averaged Magnetic Field Strength Over 4 Solar Activity Cycles |B| observed at any time is balance between flux added by CMEs and flux removed via reconnection close to the Sun.

There appears to be a floor to |B| of about 4.5 nT.

Commonly Observed Ionization States in the Solar Wind

He 2+ C 5+ , C 6+ O 6+ to O 8+ Si 7+ to Si 10+ Fe 8+ to Fe 14+ 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.

Ionization state temperatures reflect electron temperatures in the solar corona where the ionization states freeze in and are typically 1.4 - 1.6 x 10 6 ˚K in the low-speed wind and 1.0 - 1.2 x 10 6 ˚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 CMEs, reflecting the unusual coronal origins of those events.

Turbulence and Alfvenic Fluctuations in the Solar Wind

64-s data 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 velocity and magnetic field vectors). The Alfvenic fluctuations are probably remnants of waves and turbulence that heat and accelerate the solar wind. Fluctuation amplitudes decrease with increasing heliocentric distance; their dissipation heats the wind far from the Sun.

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 recent 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 10 8 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 Origin of 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