such as suprathermal electrons, ubiquitous in the solar wind, could also contribute to the acceleration of the wind by imposing an electric field on the ions, extracting them out of the corona to high speeds (Pierrard & Pieters, 2014). It is, however, very challenging to deal with these different types of processes altogether in a unified view, as they work on very different scales.
At solar minimum, the quasi‐dipolar distribution of coronal magnetic field lines extends into two magnetic hemispheres of opposite polarities in the interplanetary medium. These hemispheres are separated by a neutral line called the heliospheric current sheet (HCS). A schematic of the global structure of the HCS is shown in Figure 1.6. The HCS is typically identified in situ as an abrupt change in the magnetic field direction and a 180° switch in the pitch angle of suprathermal electrons (above tens of eV). It is also commonly called a sector boundary (Liu et al., 2014). The HCS is usually surrounded by a high‐density region called the Heliospheric Plasma Sheet (HPS) that may correspond to the heliospheric extension of coronal streamer rays (Winterhalter et al., 1994). The sector boundary and the magnetic field reversal are, at times, not collocated in situ (Owens et al., 2013). This is likely indicative of the dynamic effects, already mentioned, occurring in the solar corona near the source of sector boundaries leading to a local folding of the magnetic field that is not necessarily associated with the large‐scale HCS (Owens et al., 2013). The HCS exerts only small latitudinal displacements during solar minimum, but can extend over a broad latitudinal range at solar maximum when solar active regions are more numerous at low latitudes. For weak to moderate solar activity, the topological shape of the HCS resembles that of a ballerina skirt, as illustrated in Figure 1.6.
The expansion of the solar wind plasma combined with the effect of solar rotation shapes the interplanetary magnetic field into Archimedean spirals rooted at the Sun with a radial position r( θ )=k θ /(2 π ), where r is the radial coordinate, θ the azimuthal angle, and k is the product of synoptic period of the Sun and the solar wind speed (Parker, 1958). These spirals are known as Parker spirals and, according to Maxwell’s equations, plasma elements situated on a same Parker spiral share a common source region at the Sun. The measured dependence between the magnetic field angle with respect to the radial direction and solar wind speed agrees very well with the simple Parker spiral model. This is true in the Helios spacecraft measurements of the inner heliosphere (Bruno & Bavassano, 1997), of near 1 AU measurements taken over several decades (Borovsky, 2010), and further out along the Pioneer and Voyager orbits in the more distant outer heliosphere (Burlaga & Ness, 1993). Measurements of magnetic fields at high latitudes have been limited to those made by the Ulysses spacecraft. These measurements also confirm that the angle of the interplanetary magnetic field to the radial direction also agrees well with the Parker spiral model (Forsyth et al., 2002). Deviation from the average spiral orientations can be due to dynamic processes that occur near the source region of the solar wind or the onset and development of turbulent flows (Fisk, 1996). Both of these are discussed at length in the later sections of this chapter; they are likely important to place constraints on theories of coronal heating and solar wind formation.
Figure 1.6 Configuration in the inner heliosphere of a “ballerina skirt” heliospheric current sheet extending above the streamer belt near solar minimum (for A > 0 solar magnetic field polarity, that is, outward field at the north pole), which lies ahead of a high‐speed stream (drawn truncated at high latitudes) from an equatorward extension of a northern polar coronal hole. The dark shaded region is the interaction region.
(Source: Image reproduced with permission from Schwenn, 1990, © 1990, Springer.)
1.3.2. Composition of the Solar Winds
Figure 1.7 presents solar wind parameters during the passage of a high‐speed stream bound by two intervals of the slow wind. In situ measurements at 1 AU and beyond are not ideal to study the origin of the solar winds because the latter evolve greatly between the Sun and 1 AU. For example, the fast and slow winds interact during transit to 1 AU, and the distribution functions are strongly affected by a number of processes during transport from the Sun as discussed later in this chapter. In essence, the solar wind undergoes a number of processes that irreversibly erase information on their source regions during transit. Wind properties that remain unchanged during transit to 1AU are the abundance and ionization state of minor ions. They provide crucial information on the origins of these winds.
As an element rises in the solar atmosphere from the chromosphere to the hot corona, its ionization level increases due to radiative and collisional processes. This carries on until it reaches a height where radiative ionization is too weak and collisions are rare. Beyond that height, the charge state is said to be “frozen” and remains unchanged until it is measured in situ in the interplanetary medium. A measure of the charge state of heavy ions therefore provides unique information on the temperature at the source regions of the solar winds (Geiss et al., 1995). The charge state of an ion species will increase with temperature at the collisional coronal base of a solar wind flux tube. Solar wind measurements have shown that the carbon and oxygen ions have lower charge states (07+/06+ or C5+/C6+) in the fast than in the slow wind (Geiss et al., 1995; Kasper et al., 2012). This is clearly seen in Figure 1.7 for the 07+/06+ ratio.
Figure 1.7 Typical time profiles of solar wind parameters for a selected solar wind interval. The SW parameters are (from top down) proton speed (vp), O+7/O+6 density ratio, average charge of Fe (QFe), Fe/O relative to its photospheric value (0.0646; Asplund et al., 2009, related to the first ionization potential, FIP, bias), proton density (np), and the magnitude of the associated interplanetary magnetic field (—B—). The selected SW interval is between the two vertical dashed lines.
(Source: Image taken with permission from Ko et al., 2014. © 2014, IOP Publishing.)
Electron temperatures prevailing in the source region of the fast solar wind are therefore lower than temperatures near the source region of the slow wind (temperature smaller than 1.2 MK). Near 1 AU, where ion temperatures of the fast and slow solar winds are higher than temperatures estimated from charge state ratios. This suggests that sustained heating processes occur in the fast and slow solar wind in the corona well above the heights where ionization states are frozen (Lopez et al., 1986).
Sharp O7+/O6+ boundaries occur at both leading and trailing edges of high‐speed streams, suggesting rapid transitions between different coronal sources (Borovsky & Denton, 2016; Burton et al., 1999). High‐cadence measurements of the ionization state of heavy ions in the slow wind have recently revealed a high degree of variability, changing by an order of magnitude inside density structures (Kepko et
