The solar wind is created by the pressure difference between the solar corona and interplanetary space. This pressure difference drives plasma radially outward. If the corona was spherically symmetric in density, temperature, and pressure, the equatorial solar wind would look like the picture in Figure 1.4, where each spiral is tethered to the rotating Sun and the speed at which the plasma moves outward has angular symmetry. In this simplest case, the spiral streamlines would
Figure 1.26: Sketch of the magnetic field lines in the corona during a solar minimum and solar maximum. Adapted from Forsyth (2001).
not intersect. The real corona however is not spherically symmetric. Instead, the solar magnetic field evolves on both short and long timescales. At the beginning of a solar cycle, the magnetic field in the corona is roughly dipolar, with a slight tilt. The left panel in Figure 1.26 shows the idealized solar magnetic field structure at solar minimum. During this time, the polar regions have field lines which are open and only tethered to the Sun on one end, while equatorial regions have closed field lines. Regions of open field lines are referred to as coronal holes and the loop-like structures created by closed magnetic field lines are called helmet streamers. The solar wind emanating from open field lines is much faster than the solar wind emanating from regions with helmet streamers. The plasma inside the Sun is not static. As the solar cycle progresses, the buoyant rise of active regions impacts the solar magnetic field. At solar maximum the solar magnetic field looks more like the right panel in Figure 1.26.
The shape of the solar wind streamlines in Figure 1.4 are speed dependent. If at the equator a coronal hole follows a helmet streamer, the faster solar wind from the hole will catch up to the slower solar from the streamer. Since the slow and fast streams originate from different positions and at different times from the Sun, their frozen-in magnetic fields are different and the two streams cannot interpenetrate. This leads to a compression of the solar wind plasma between the faster and
slower streams, a steadily rising speed between the two, and rarefaction on the trailing edge of the fast stream. These features are shown in Figure 1.27. The Sun rotates roughly once every 27 days, and if the flow patterns emanating from the Sun are stable over this period, an observer at a fixed distance from the Sun will see this fast/slow stream interaction with the same periodicity. These stream interactions are called co-rotating interaction regions (CIRs). The temporal variability in the solar-wind structures can be on a shorter timescale than the solar rotation period and the periodic stream interaction might not be observed. One-off stream interactions without periodicity are referred to asstream interaction regions(SIRs) and include possibly localized stream interactions.
Figure 1.27: Cartoon co-rotating interaction region. Adapted from Pizzo (1978).
Figure 1.28 shows the typical solar wind ob- served by the Wind spacecraft at L1 (Lagrangian
point of a satellite balanced between the forces of gravity between the Sun and the Earth, located∼.01 AU from the Earth to the Sun). In stream inter- actions, the boundary separating the originally fast stream and the slow stream is referred to as the stream interface. The stream interface often features
an abrupt drop in particle density and a simultaneous rise in proton temperature as the faster stream is often hotter and more tenuous than the slower stream [Gosling et al. (1978b)]. In Figure 1.28, the stream interface is denoted by the black vertical line. TheWinddata shows that leading up to the stream interface the magnetic field strength, density, and solar wind speed increase due to the compression. Following the stream interface, the speed and temperature increase as the satellite observes plasma originating from the coronal hole, but the density and magnetic field strength fall off after the compression. The total pressure (combined magnetic pressure and thermal pressures for all ions/electrons) reveals a pyramidal shape in the observations by a stationary satellite, and often the peak in total pressure can be used to identify the stream interface [Jian et al. (2006c)]. Zero crossings of the azimuthal flow angle of the solar wind can be used to identify a stream in- terface [McPherron et al. (2009)]. Statistical studies have shown that over a solar cycle,∼25%of
Figure 1.28: Observation of a corotating interaction region. Data from theWindspacecraft. Verti- cal line denotes the stream interface.
SIRs have either a leading forward, or trailing reverse shock at 1 AU, but that further away from the Sun that number increases to > 80%by Jupiter’s orbit [Jian et al. (2008, 2006c)].
CIRs are important to the magnetosphere as they have high speed solar wind and can posses a strong magnetic field which can drive dayside reconnection when they reach the magnetic field of the Earth. The zˆcomponent of the magnetic field in CIRs is quasi-random in its orientation, and can often have high degrees of Alfvénic fluctuations. Short periods where the zˆcomponent is negative will drive dayside reconnection and enhancements to the convection electric field, as well as ionospheric outflow from the polar regions to the plasma sheet. Thus, CIRs can generate geomagnetic storms. CIRs can often last for multiple days, but often are not the strongest storms, as there is no preferential orientation of thezˆcomponent of the magnetic field, so it is not guaranteed to drive steady convection for long periods of time [Gonzalez et al. (1994)]. Typically, around 30–40 CIR events reach L1 each year, and they are most prevalent around solar minimum [Jian
et al. (2008, 2006c)].