It was not until the Skylab era in 1973/74 when so-called coronal holes were discovered to be the sources of long-lived solar wind high-speed streams (Krieger et al., 1973). Coronal holes are usually located above inactive parts of the Sun, where “open” magnetic field lines prevail, e.g., at the polar caps around activity minima. In contrast, the more active near-equatorial regions on the Sun are most often associated with closed magnetic structures, such as bipolar loop systems and helmet streamers on top. From here, the more turbulent slow solar wind emerges. The embedded density fluctuations allow visualization of this type of solar wind. The images of Figure 6 (taken with the most sensible coronagraph LASCO C3 on SOHO) reveal that it is ejected continually near the equatorial plane, with the appearance of cigarette smoke. Sometimes the patterns in the out-flowing clouds can be tracked like “leaves in the wind” such that Sheeley Jr et al. (1997) could even determine the speed profile. Unfortunately, the high-speed streams above the poles remain invisible, they do not have sufficiently strong density fluctuations.
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It is important to note that both: the coronal holes as well as their offspring, the high-speed solar wind streams are representatives of the inactive or “quiet” Sun. Thus, the only state of the solar wind that may deserve the label “quiet” is the high-speed wind, rather than the more variable slow wind from above active regions. This perception, first described by Feldman et al. (1976) and Bame et al. (1977), caused a major paradigm change. No longer could the slow wind be considered the “quiet” or “ground state” type although it would fit much better to the famous model of a thermally driven solar wind as derived by Parker (1958).
Figure 7 exhibits the two states of the near-minimum corona in 1998 and the associated solar wind streams rather nicely. The parts of this combined image were taken almost simultaneously from the EIT and LASCO telescopes on SOHO. This figure shows the two states by their different brightness in the corona. Note also how well separated from each other they are, both in the low corona as well as in the extended corona, i.e., in the solar wind.
In fact, the existence of sharp boundaries between solar wind streams (in longitude as well as in latitude) had already been noticed by Rosenbauer et al. (1977) and Schwenn et al. (1978) on the basis of in situ measurements from the Helios solar probes that went as close as 0.3 AU to the Sun. These two basic types of quasi-steady solar wind differ markedly in their main properties and by the location and magnetic topology of their sources in the corona, thus probably in their acceleration mechanism.
The main characteristics of the two basic states as determined by Schwenn (1990) from measurements in the plane of the ecliptic are listed in Table 1.
|Low speed wind (LSM)||Fast wind (HSS)|
|Flow speed vp||250 – 400 km s−1||400 – 800 km s−1|
|Proton density np||10.7 cm−3||3.0 cm−3|
|Proton flux density np vp||3.7 × 108 cm−2 s−1||2.0 × 108 cm−2 s−1|
|Proton temperature Tp||3.4 × 104 K||2.3 × 105 K|
|Electron temperature Te||1.3 × 105 K||1 × 105 K|
|Momentum flux density||2.12 × 108 dyne cm−2||2.26 × 108 dyne cm−2|
|Total energy flux density||1.55 erg cm−2 s−1||1.43 erg cm−2 s−1|
|Helium content np∕nHe||2.5%, variable||3.6%, stationary|
The main differences between the two states are, apart from the speed itself, the particle densities, the mass flux, and the helium content. However, note the striking similarities in the flux densities of energy and momentum. The latter is equivalent to the ram pressure the solar wind is exerting on an obstacle like the Earth. It appears to be invariant with the flow state within 7%. The same result was found to apply for the latitudinal stream structure up to at least 30∘ on heliographic latitude (Bruno et al., 1986). That means that size and shape of the Earth’s magnetosphere are not affected much by the solar wind stream structure. That aspect should be kept in mind in the space weather context.
The total energy flux density which is also invariant is the sum of two main components: the kinetic energy flux and the potential energy flux (basically the work done for moving the solar wind out of the Sun’s gravitational potential). Although these components differ strongly for the two states, their sum is about the same. There is no explanation yet for these strange invariances. We suspect that a crucial key to understanding the solar wind phenomenon is hidden here.
In addition to these two basic states, the slow solar wind filling most of the heliosphere during high solar activity can be considered a third category. It emerges above active regions distributed over large parts of the Sun, far from the heliospheric current sheet, and in a highly turbulent state. It differs in some aspects from the minimum type of slow solar wind. Finally, we regard the plasma expelled from the Sun during huge coronal mass ejections as a category on its own, because of some fundamental differences to be described later.
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