The solar wind

2.1 The solar wind

Before describing how solar-like winds are detected around other stars, it is worthwhile to briefly review what is known about the solar wind and how we study its properties. The solar wind was first detected through its role in the formation of aurorae and the creation of comet tails. As far back as 1896, Kristian Birkeland proposed that aurorae were due to particles emanating from the Sun (see review by Stern, 1989 ), while Biermann (1951 ) first described how “corpuscular radiation” from the Sun was responsible for the plasma tails of comets. Today almost everything we know about the solar wind comes from in situ measurements of its properties from satellites. These measurements date back to the Soviet Luna missions in 1959 (Gringauz et al., 1962 ) and NASA’s Mariner 2 mission in 1962 (Neugebauer and Snyder, 1962 ). Numerous other spacecraft have participated in studying the solar wind since then, but of particular note are the venerable Voyager 1 and Voyager 2 satellites, which have returned data on the solar wind from 1977 through the present day (see Lazarus and McNutt Jr, 1990

). More recently, the Ulysses spacecraft, launched in 1990 and still operating, has provided a first look at the solar wind outside of the ecliptic plane (McComas et al., 2000

Figure 1:

The solar wind velocity (red/blue line) and density (green line) observed by Ulysses as a function of ecliptic latitude (McComas et al., 2000

). During solar minimum conditions, high latitudes are dominated by high speed, low density wind, while low latitudes see mostly lower speed wind with higher densities.

Within the ecliptic plane the solar wind is dominated by low speed streams with typical velocities, proton densities, and temperatures (at 1 AU) of $V = 400 km s− 1$ , $n(H+ ) = 5cm −3$ , and $T = 105K$ , respectively, although high speed streams with lower densities and $V ≈ 800km s−1$ are not uncommon (see Feldman et al., 1977 ). At the maximum of the 11-year solar activity cycle, similar solar wind behavior is seen at almost all latitudes. However, Figure 1 shows that at solar minimum the wind above $∘ 30$ ecliptic latitude is uniformly high speed, low density wind with $−1 V ≈ 800 km s$ (McComas et al., 2000 , 2002 ). This high speed wind originates from coronal holes, which are particularly prominent on the Sun during solar minimum conditions. These solar wind data imply a total mass loss rate for the Sun of $M˙ ≈ 2 × 10 −14M yr− 1 ⊙ ⊙$ . Although thermal temperatures for the wind at 1 AU are of order $5 T = 10 K$ , densities are low enough that the wind cannot equilibrate to this temperature, and the ionization state of the wind is actually frozen in at coronal temperatures closer to $6 T = 10 K$ . Hydrogen is fully ionized, meaning that protons are the dominant constituents of the wind by mass.

Perhaps the most fundamental question to ask about the solar wind is why it exists. Addressing this question is also necessary to assess whether similar winds should exist around other stars. Even before satellites proved the existence of a more-or-less steady solar wind, Parker (1958 ) predicted that such a wind should be present due to the existence of the $106 K$ corona surrounding the Sun. In Parker’s model, the solar wind exists because of thermal expansion from the hot corona. The predictions of this simple model agree remarkably well with the observed properties of the low speed wind that dominates in the ecliptic plane, although additional wind acceleration mechanisms invoking MHD waves have been proposed to explain the high speed streams (see MacGregor and Charbonneau, 1994 ; Cranmer, 2002 ). Thus, any star that has a hot corona analogous to that of the Sun should also have a wind analogous to that of the Sun. Observations with X-ray satellites such as Einstein and ROSAT clearly demonstrate that coronae are a ubiquitous phenomenon among cool main sequence stars (see Schmitt, 1997 ; Hünsch et al., 1999 ), so solar-like winds should be present around all solar-like stars. However, that does not mean that they are easy to detect (see Section 1).