discussed later in the chapter.
The higher charge state of heavy ions measured in the slow wind necessarily results from higher temperatures and densities at the coronal base of open field lines channeling the slow wind. Numerical models of the solar coronal plasma and magnetic field have been used to study the origin(s) of the slow wind. They reveal that magnetic fields are generally stronger at the base of flux tubes channeling the slow wind (Wang et al., 2009). There is a clear statistical correlation between coronal regions threaded by strong magnetic fields in the form of loops or open magnetic fields and high‐plasma temperatures (Schrijver et al., 2004). Therefore, one possible interpretation for the high charge‐state ratios measured in the slow wind resides in strong local heating near its source region due to the strong magnetic fields (Wang et al., 2009). Such a strong heating at the coronal base would lead to a strong heat flux conducted down to the chromosphere and enhanced densities in the slow wind. In contrast, less evaporation is likely to occur near the cooler source region of the fast wind, leading to a more tenuous fast wind. This interpretation can explain the fairly constant mass flux measured in situ in the fast and slow solar winds.
Figure 1.8 Element abundances as a function of the first ionization potential (FIP) in the average slow solar wind, fast solar wind. Abundances are given relative to oxygen and are normalized to photospheric abundances.
(Source: This figure was taken from Geiss, 1998 with permission from SSR. © 1998, Springer Nature.)
An alternative theory for the origin of the high charge states measured in the slow wind suggests that plasma initially confined to coronal loops is released along open magnetic field lines (Fisk et al., 1998; Schwadron et al., 1999). Loop plasma would naturally be pushed to higher ionization states because loops host typically hotter and denser plasma than open fields of coronal holes (Schwadron et al., 1999). For this mechanism to work, loop plasma must find a way to be transferred to the open field, which would require magnetic reconnection to occur continually.
The fast and slow solar winds also differ in their abundance of elements with low First Ionisation Potential (FIP), that is, the minimal energy required to ionize an atom. This first ionization occurs in the chromosphere for most elements. Overall, the corona and the solar winds exhibit higher abundances of elements with low FIP than known photospheric abundances. The slow wind exhibits even greater abundances of low‐FIP elements such as Si, Fe, and Mg than the fast solar wind (von Steiger et al., 2000; Zurbuchen et al., 1999) as seen in the fourth panel of Figure 1.7 and in Figure 1.8. Such high abundances of low‐FIP elements are not measured remotely in coronal holes but are more typical of closed coronal loops (Widing & Feldman, 2001). This supports the idea that the slow wind is supplied by plasma initially trapped on coronal loops and that are subsequently released in the wind (Fisk et al., 1998). In contrast, the fast wind carries fewer elements with low FIP (von Steiger et al., 2000). The composition of the fast solar wind agrees roughly with that measured in coronal holes.
Figure 1.9 Left: The abundance of helium (computed as 100 times the density ratio of hydrogen) during 17 years of solar wind measurements by the Advanced Composition Explorer (ACE). The color coding indicates the solar wind speed. The monthly smoothed sunspot number is plotted as a black curve. Right: The abundance of helium relative to oxygen, measured by ACE/SWICS. The same trend of depletion with wind speed as for He/H is seen, but the solar cycle dependence is less pronounced.
(Source: Images reproduced with permission from Kasper et al., 2012 and Rakowski & Laming, 2012. © 2012, IOP Publishing.)
The fast and slow winds can also differ in their abundance of elements with higher FIP such as helium. As shown in the left‐hand panel of Figure 1.9 taken from Kasper et al. (2012), the helium abundance of the fast wind remains steady (4–6%) throughout the solar cycle. In contrast, the helium abundance in the slow wind varies from around 1% at solar minima to around 4–6% at solar maxima (Aellig et al., 2001; Kasper et al., 2012). Therefore, the slow wind carries less helium than the fast wind at solar minima, but both present similar abundances at solar maxima.
There is currently no accepted mechanism for the regulation of heavy ion abundances in the solar wind. All proposed mechanisms must necessarily occur where elements are first ionized in the chromosphere and fractionate elements in a mass‐independent manner. Several mechanisms have been invoked in this region to modulate the transfer of heavy ions from the chromosphere to the corona. They include, for example, the effect of Coulomb collisions in high‐temperature gradients (Bø et al., 2013) and the effect of MHD waves via a ponderomotive force (Laming, 2009, 2015). The latter mechanism, in particular, is able to explain a broad range of composition measurements in the fast and slow solar winds for reasonable conditions in the solar atmosphere.
1.3.3. Solar Wind Interaction Regions
The rotation of the Sun and its associated solar wind sources such as coronal holes can result in the radial alignment of the fast and slow solar wind. A kinematic gathering of the interplanetary plasma when the fast wind catches up the slow stream, inducing an increase in plasma density near the stream interface. When the solar wind behind is much faster than the solar wind ahead, it can compress the slow wind before it reaches 1AU (as seen in Figure 1.7). When a compression region is measured over at least two consecutive solar rotations it is called a Corotating Interaction Region (CIR), otherwise one refers to simply a Stream Interaction Region (SIR; L. Jian et al., 2006). The interaction between fast and slow solar wind that creates a CIR is illustrated in Figure 1.10. Detailed descriptions of CIRs can be found in Gosling and Pizzo, 1999, and in Pizzo, 1982. As seen in the left‐hand schematic of Figure 1.10, the interaction region forms, like the background solar wind, an Archimedian spiral rooted at the Sun. In three dimensions, the stream interface can be titled as it responds to the tilt of the solar magnetic dipole. Figure 1.10 illustrates how plasma flows are typically deflected inside CIRs in latitude as well as in longitude (Pizzo, 1982).
CIRs and SIRs can be more complex structures than spirals; they are composed of a multitude of smaller interaction regions that merge during the wind propagation to 1 AU. This is a consequence of the variability of the slow wind released by helmet streamers (as discussed in Section 1.2.2). Strong dynamic effects associated with the rise in total pressure inside the SIRs and CIRs are most apparent beyond 1 AU, where forward/reverse shock pairs (Lee, 2000) can bound the interaction region.
