the larger plasmoids (Kepko et al., 2016; Viall et al., 2010; Viall & Vourlidas, 2015). DeForest et al. (2018) used deep‐field, high‐cadence coronagraph observations to show that there is still more substructure (shown in Figure 1.4)—both time dynamic on scales smaller than 90 min, and time stationary, filamentary streamer structures down to the resolution limit. The time‐dynamic characteristic scale sizes likely have different formation causes, such as inherent time scales due to the characteristics of coronal heating (Endeve et al., 2004), or waves (Pylaev et al., 2017).
Observations of the corona therefore show that highly dynamic streamers host transient processes associated with the likely continual emergence, redistribution, and removal of solar magnetic flux directly influencing the properties of the magnetic fields and particles of the solar wind (Owens et al., 2013; Sanchez‐Diaz et al., 2016). Magnetic reconnection occurring at the boundary of coronal holes was suggested to occur in the 1980s to explain the rigid rotation of coronal holes (Wang et al., 1988). Theoretical considerations (Fisk, 1996) and sophisticated numerical MHD simulations have highlighted the inevitable, complex evolution that takes place all along the corona’s open‐closed field boundaries (Antiochos et al., 2007; Linker et al., 2011; Lionello et al., 2005; Titov et al., 2009; Titov et al., 2011). The activity hosted by helmet streamers are the clearest example, with many different scales of plasma and magnetic field signatures generated by magnetic reconnection (Higginson et al., 2017; Higginson & Lynch, 2018; Lionello et al., 2005; Török et al., 2009).
Heliospheric imaging has also shown that not all structures at MHD scales measured in the solar wind form in the corona. DeForest, Matthaeus, Viall, and Cranmer (2016) showed turbulent density fluctuations setting in around 30 solar radii away from the Sun. These fluctuations coexist with outflowing helmet streamer plasmoids and are advected with the solar wind.
1.3. MEASUREMENTS OF THE SOLAR WIND IN THE INNER HELIOSPHERE
1.3.1. Bulk Properties and Large‐Scale Structures
The properties of the solar wind escaping the corona change throughout the solar cycle. This results from the evolving coronal magnetic topology that responds to the emergence and evolution of photospheric magnetic fields. This evolution alters both the magnetization and the bulk properties of the wind at different heliocentric latitudes. Spacecraft have measured in situ the properties of solar wind magnetic fields and particles at different locations in the heliosphere and over several decades. Figure 1.5 displays three dial plots showing the distribution of solar wind speeds with latitude measured by the Ulysses spacecraft during its three polar orbits. The sunspot number plotted below the dial plots is low in the left‐hand and right‐hand plots, marking the occurrence of solar minima (McComas et al., 2008). At these times, the Sun’s magnetic field is quasi‐dipolar, and the ambient solar wind is very clearly structured in at least two types of plasma flows. A fast wind is measured at high latitudes above coronal holes, and a more complex slow wind is measured at the low latitudes of solar streamers. As the solar cycle advances toward the activity maximum (middle panel), the polar coronal holes can disappear temporarily and the magnetic field evolves toward a non‐dipolar structure. In response to this process, the bulk speed loses the ordered latitudinal structure shown in Figure 1.5 (left). At solar maximum, the large‐scale spatial separation between the slow and fast solar wind is less clear as shown in the middle dial plot. Global numerical models of the solar wind show the clear link between these changes in the coronal magnetic topology and wind streams (Linker et al., 2011; Oran et al., 2013; Pinto et al., 2011; Pinto et al., 2016; van der Holst et al., 2014).
Figure 1.5 (a–c) Polar plots of the solar wind speed, colored by IMF polarity for Ulysses’ three polar orbits to indicate measured magnetic polarity. (d) Contemporaneous values for the smoothed sunspot number (black) and heliospheric current sheet tilt (red), lined up to match Figures 1.1a–c. In Figures 1.1a–c, the solar wind speed is plotted over characteristic solar images for solar minimum for cycle 22 (17 August 1996), solar maximum for cycle 23 (7 December 2000), and solar minimum for cycle 23 (28 March 2006). From the center out, we blend images from the Solar and Heliospheric Observatory (SoHO) Extreme ultraviolet Imaging Telescope (Fe XII at 1950 nm), the Mauna Loa K coronagraph (700–950 nm), and the SoHO C2 white‐light coronagraph.
(Source: Image reproduced with permission from McComas et al., 2008, © 2013 John Wiley & Sons.)
The first models capable of describing the general properties of the solar wind assumed a thermally driven flow. These include Parker’s original (1958) theory that assumed a constant coronal temperature as well as subsequent fluid models that allowed for thermal stratification and inhomogeneities. The latest fluid models account for more detailed energy injection and transport mechanisms (Linker et al., 2011; Lionello et al., 2014; Oran et al., 2013; Pinto & Rouillard, 2017; van der Holst et al., 2014). Fluid models are not able, for known coronal temperatures, to explain the high speeds of the fast solar wind without including additional physical processes than just the effect of the thermal pressure gradient. These processes could involve Alfvén waves with their induced turbulent wave pressure and Reynolds stresses that would contribute to further accelerate the solar wind (Chandran, 2018; Cranmer et al., 1999; Lionello et al., 2014; Oran et al., 2013). Kinetic solar wind models suggest that heated particles 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