Presentation on theme: "The Solar Wind J. T. Gosling LASP, University of Colorado June 11, 2009."— Presentation transcript:
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 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.
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 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 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 is solar wind speed, B is magnetic field strength, A(He) is He ++ /H + ratio, T p is proton temperature, T e is electron temperature, T is alpha particle temperature, C s is sound speed, C A is Alfven speed. 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 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 (B r positive or negative) throughout. Reversals in B r 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. 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 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. F R
Evolution of Stream Structure with Heliocentric Distance (continued) 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. F R
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 Stream amplitudes are strongly damped in the outer heliosphere because of the interactions between high and low-speed streams. Voyager 2 data obtained at ~18 AU
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 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. F R
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 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 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 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 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