a very small region of the heliospheric boundary.
Birkeland (1908) argued very early that a corpuscular emission from sunspots consisting of relativistic electrons must impact Earth’s magnetic field and be deflected to the polar regions to create the aurora. For several decades, it was realized that particles could be emitted from the Sun during flares, but it was generally thought that the space around Earth was mostly empty or perhaps traversed by occasional streams of particles from the Sun (Chapman & Ferraro, 1931). The prevailing view was that the solar corona consists of a hot gas (possibly extending to 1 AU), in thermal and hydrostatic equilibrium, pulled back by the solar gravitational field (Chapman & Zirin, 1957). Detailed observational studies of comets by Biermann (1951) showed that a subset of their tails cannot be accelerated by radiation pressure alone but may also respond to material flowing out from the Sun’s atmosphere. He suggested that the passage of solar particles at the comet formed an ion tail and that these particles must have a very high speed relative to the comet in order to align the tail in the Sun’s direction. Parker (1958) built on these observations and realized that the high temperature of the corona can provide enough energy to force coronal plasma to accelerate from subsonic to supersonic speeds. He demonstrated that the hydrostatic approach predicted too high kinetic pressure at infinity and that a continuous radial expansion of solar gas must act to reduce the coronal pressure. This was the first theory describing the continual expansion of what we now call the solar wind.
In this model, a dominant force affecting coronal particles and pushing them outward is induced by the thermal pressure gradient in the corona. Parker’s original model assumed an isothermal corona, but subsequent models allowing for a varying temperature with radial distance confirmed that a supersonic wind can also form under such conditions. The presence of an outflowing supersonic wind ranging from 300 to 800 km/s was confirmed by plasma measurements from the Luna 1, 2, 3; Venera 1; and Mariner 2 (e.g., Neugebauer & Snyder, 1962) and the numerous subsequent solar wind dedicated missions (e.g., Hundhausen & Gosling, 1976; Marsch et al., 1982). The Parker model provided an acceleration from subsonic speed in the coronal region to supersonic speeds of around 300–400 km/s typical of the slow solar wind. His model was unable, however, to explain fast solar wind speeds of 700–800 km/s without assuming unrealistic coronal hole temperatures in excess of 2 × 106K. As we shall show in this chapter, coronal and solar wind measurements suggest that additional physical mechanisms must be accounted for to produce a fast solar wind.
The solar wind has been measured in situ for several decades near 1 AU. Typical solar wind speeds near 1 AU range from 300 to 800 km/s, proton temperatures take values between 105 K in the slow wind to about 2 × 105 in the fast wind. There is now no doubt that the fast solar wind measured in situ originates near the center of coronal holes at the Sun (see Chapters 2 and 3). Coronal holes are cooler regions of the solar atmosphere that exhibit drops in Extreme UltraViolet (EUV) emissions. The cooler temperatures result from the significant escape of heat out in the solar wind. The fast solar wind streams out along open magnetic fields connecting regions deep inside coronal holes to the interplanetary medium. The origin of the slow wind is more complex and likely consists of multiple source components: these include transient releases from helmet streamers, plasma accelerating on rapidly expanding magnetic fields rooted at the boundary of coronal holes (through processes likely similar to that of the fast wind), and/or continual plasma exchanges between loops and open magnetic fields. We will describe these different components in the following sections.
1.2. OBSERVATIONS OF THE NASCENT SOLAR WIND
1.2.1. Remote‐Sensing Observations of Coronal Heating and the Solar Wind
We begin our story on the solar wind with its formation in the solar corona; we first present remote‐sensing observations that have provided important information on the conditions in which the winds are produced.
It is hard to observe the coronal source regions of the fast and slow solar winds in white‐light images obtained routinely by coronagraphs. This is because the fast solar wind originates in very tenuous coronal holes, whereas the slow solar wind emerges from the vicinity of dense streamer loops that completely dominate coronal brightness. In contrast, spectroscopic observations provide more detailed information about the temperatures, flow velocities, and wave properties during the formation of the winds near their sources.
The Solar and Heliospheric Observatory (SoHO; Domingo et al., 1995) revolutionized how we observe the corona and the nascent solar wind. In particular, the Ultraviolet Coronagraph Spectrometer (UVCS/SoHO; Kohl et al., 1995) revealed that heavy ions such as O5+ and Mg9+ are heated hundreds of times more strongly than protons and electrons, and have very anisotropic kinetic temperatures (Antonucci et al., 2000; Kohl et al., 1997, 1998), meaning that temperatures measured in the direction perpendicular to the magnetic field are often much larger than those parallel to the field (see Chapter 2). A comparison between different ion and electron temperatures derived from coronal observations and solar wind measurements are shown in Figure 1.1. The measured temperatures provide support to coronal heating mechanisms involving collisionless wave–particle resonances with frequencies of 10 Hz to 10 kHz. These waves could be damped to heat preferentially heavy ions (Cranmer et al., 1999; Tu & Marsch, 1997). It is not clear yet how such hypothetical waves can be generated from lower‐frequency Alfvén waves typically emitted at minute periods. Some proposed mechanisms involve the turbulent cascade of magneto‐plasma fluctuations from low to high frequencies (Hollweg, 2002). Then kinetic processes must occur to damp the small‐scale fluctuations that have spectrally cascaded to oblique wavevectors. Proposed mechanisms include ion‐cyclotron and Landau damping (Leamon et al., 1998). Nonlinear processes such as beam instabilities or mode conversion and damping have also been proposed.
Figure 1.1 Radial evolution of solar wind temperatures from the corona to 1 AU. Indirect estimates of coronal temperatures are derived for three species (electrons, oxygen, and hydrogen) from an “empirical coronal model” that exploits SUMER (1–1.2R⊙) and UVCS (1.5–4R⊙) line widths compared with direct in situ measurements at distances greater than 60R⊙) in the high‐speed wind (Cranmer et al., 1999). The in situ data were assembled from Helios, IMP, Ulysses, and Voyager particle data, and double sets of curves denote rough lower and upper bounds on representative fast‐wind values. The upper and lower limits correspond to extreme values of an assumed non‐thermal component to the line broadening that is attributed to unresolved MHD wave motions along the line of sight. Temperatures of electrons (solid black), hydrogen (dotted), and oxygen (dashed) are shown. Oxygen ions correspond to O5+ in the corona but O6+ in the far solar wind, and coronal temperatures of neutral hydrogen are here compared with proton temperatures in the solar wind. Figure taken with permission from (Cranmer et al., 1999).
Emission lines, and more specifically the dimming of certain lines measured via coronal spectroscopy, have been used to infer the outflow speed and density of the forming fast solar wind. This effect is most important for spectral lines in the UV range that have a significant component due to ion excitation via resonant scattering of chromospheric emission (followed by spontaneous emission). For some lines, such as O VI doublet 103.2 nm and 103.8 nm lines, and of the HI Lyman‐ α 121.6 nm and Lyman‐ β 102.5 nm lines, the collisional and radiative components can be separated (Marocchi et al., 2001). This separation can be used to investigate the dimming of the radiative component, which is due to the radial expansion of the emitting coronal atoms. This Doppler shifts the (narrow) exciting chromospheric profile with respect to the (broad) atomic absorption profile. As a result, the UV intensity of the resonantly scattered
