Overview and Current Knowledge of the Solar Wind and the Corona

Fast and Slow Solar Wind. 

Ulysses, with its near polar 1.4 by 5.4 AU orbit, has given us a graphic picture (at the right) that solar wind comes in two states: an irregular slow wind with typical speeds of 400 km/s and a smooth fast wind with a speed of ~750 km/s. This is the "bimodality" of the solar wind, and is most apparent at solar minimum (when the data shown in the figure were taken). Fast wind comes from coronal holes and slow wind from the boundaries or interior of streamers. Solar Probe will encounter streamers in both 2010 and 2015 and will pass through coronal holes at 5-10 RS in 2015.

Solar wind flow speed data is displayed in the "dial plot" at the right for the Ulysses "fast latitude scan". This was the interval from September 1994 to July 1995 when Ulysses swept from 2.2 AU and 80 degrees south heliographic latitude, up to the equator at ~1.4 AU, and then on up to 80 degrees north - again at 2.2 AU. This occurred very near solar sunspot minimum when there were well defined, large polar coronal holes and a bright equatorial streamer belt - very like how the Sun appeared when the image at the beginning of the introductory page was made. 

Ulysses dial plot

Fast Wind is Steady and Simple. 

Fast wind is relatively steady (see above) and also relatively simple in composition.  The charge state distribution is characterized by a single, low freezing-in coronal temperature for each element. The elemental composition of the solar wind is least biased in the fast wind and resembles most closely the photospheric composition, and the overabundance of low FIP (first ionization potential) elements is at most weak. Fast wind is permeated by an evolving field of magnetohydrodynamic (MHD) turbulence which is presumed to be a remnant or imprint of the coronal acceleration process. 
Ulysses dial plot

Ulysses dial plot

The Slow Wind is Variable and Complicated. 

Slow wind is highly variable in speed (see above) and more complicated than fast wind in its other characteristics. The charge state distribution can no longer be characterized by a single freezing in temperature. The FIP effect is far more pronounced and the 3He/4He ratio is both higher and more variable in the slow wind compared to the fast wind. MHD turbulence in slow wind is less evolved and more intermittent than in fast wind. 

Boundary Between the Fast and Slow Wind is Sharp. 

The sharp boundary between fast and slow wind is also sharp in freezing-in temperature and FIP strength. This boundary thus must extend down to the lower corona, where the charge states freeze-in, and to the chromosphere, where the composition is established. In the above figure, the boundary is so thin that it is essentially the thickness of the line between data points defining adjacent fast and slow wind. 

The boundary between the equatorial slow wind and south polar fast wind in the figure immediately above this paragraph is a good example of this phenomenon. 

Coronal Structure and the Solar Cycle. 

The corona changes dramatically over the solar cycle, with coronal holes dominating at sunspot minimum and essentially absent at solar maximum. Streamers dominate the corona outside of coronal holes. Solar Probe will pass through the corona at both maximum and minimum to provide good data on both steamers and coronal holes.
Coronal Structure and the Solar Cycle
(a) Solar Maximum              (b) Declining Phase (plumes actually appear at all phases, but are only shown here)              (c) Solar Minimum

Characteristics of the Initial Solar Wind in Coronal Holes. 

SOHO and interplanetary scintillation results show that fast wind reaches its terminal speed by 10 Rs and that at 4 Rs it is already being accelerated. 

At 4 Rs the temperature of heavy ions is much larger than that of protons while the difference is smaller at 1 AU. The proton temperature at 4 Rs in coronal holes is two to three times higher than the electron temperature inferred from charge state measurements in the terminal wind while they differ by less than a factor of two at 1 AU. Inferred ion temperature anisotropies are enormous between 2 and 10 Rs and are believed due to an Alfvén or ion-cyclotron wave field contributing to the perpendicular temperature. A true proton temperature anisotropy exists in the 1 AU fast solar wind, but is smaller than inferred from the coronal observations. 


Plumes permeate all coronal holes, yet are invisible in the solar wind. How this variable, filamented flow becomes the uniform fast wind in unknown. Whether this is related to the source and evolution of MHD turbulence in the solar wind will be answered by Solar Probe. 

The picture on the right is a composite of SOHO/EIT, HAO Mk3 coronagraph, and SOHO/LASCO C2 showing plumes in the northern polar coronal hole in March 1996 (DeForest et al.). The plumes are seen to originate in the photosphere, diverge along with the coronal hole boundaries, and extend here to at least 6 Rs. Recent (1999) observations show plumes are sometimes detectable out to 15-30 Rs

A plume will tend to last for several days but within a plume there are brightness fluctuations on much shorter time scales. 

Solar Probe will fly directly through plumes in the polar coronal holes. 

composite of SOHO/EIT, HAO Mk3 coronagraph, and SOHO/LASCO C2 showing plumes in the northern polar coronal hole in March 1996

Characterisitcs of the Initial Solar Wind in and above Streamers. 

SOHO and other observations indicate flow speeds in and around streamers are consistent with the origin of slow wind. However, how this happens has not been determined. One problem is that the standard concept of a streamer is a magnetostatic structure which releases no wind in the steady state. Conversely, SOHO has clearly shown sporadic escapes of mass from the tops of streamers which seem to ride on a preexisting slow flow. Solar Probe will pass through the tops of streamers, precisely where this process is occurring. 

In the synthetic images below, built up from SOHO/UVCS slit spectrograms, the field of view extends approximately from 1.6 Rs out to 6.0 Rs. Again, it can bee seen that Solar  Probe will pass directly through the part of the corona most important for determining the characteristics of the initial solar wind in and above streamers. 


Properties of the Polar Photosphere. 

SOHO has hinted at some remarkable features for the polar photosphere. But, since the poles cannot be viewed effectively by SOHO (or any spacecraft confined to near the ecliptic plane), these features remain poorly defined.  These findings are: (1) the rotation rate at higher latitudes is 10% to 20% lower than expected; (2) there is some evidence for a polar vortex; (3) there is some evidence of a polar concentration of magnetic flux; (4) measurements of surface and subsurface motion indicate meridional flows that are a factor of more than two higher than previously estimated; (5) there are indications that small and large scale magnetic fields on the sun are rooted at different depths in the convection zone. These results, combined with the more general SOHO result that magnetic flux is replaced very rapidly everywhere on the surface of the Sun - approximately every 40 hours, suggest the value of a close examination of the photospheric dynamics and magnetic fields to extend our understanding of how these relate to the flow of energy into the corona. 
(doppler shift picture of solar poler regions courtesty of SOHO/SOI)
doppler shift picture of solar poler regions courtesty of SOHO/SOI

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  • Page 2: First page of the Solar Probe Science Section. 
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  • Prepared by Steve Suess (steve.suess@msfc.nasa.gov)
    3 June 1999