As stated in the previous section, the early acceleration phase of typical CMEs has mostly ceased by the time it has reached around . This indicates that over this distance the CME, which has a mass of the order of 1013 kg must be accelerated to speeds of several hundred km s–1 (and sometimes exceeding 2000 km s–1). Hence, the erupting magnetic structure that becomes the CME must have access to a mechanism providing vast amounts of energy over a relatively short time scale. Some theoretical models suggest that this involves an interaction between the erupting and the surrounding field, perhaps via runaway magnetic reconnection. The CME onset itself must involve some instability disrupting the equilibrium between the closed magnetic field in the corona and the tendency of the corona towards its natural state of expansion. We have thus far been unable to directly observe this mechanism or the instability responsible for its onset, although we can find clues via near-solar-surface phenomena that are known to be associated with the initiation of CMEs. These are mostly observed with instruments other than coronagraphs, typically imagers observing various regions of the electromagnetic spectrum: this makes the direct association between a CME and the associated phenomena difficult. In this section we review these phenomena. We also draw the reader’s attention to other recent reviews, including Webb (2002), Cliver and Hudson (2002), Gopalswamy (2004, 2010b), Kahler (2006), and Howard (2011b).
The release of the stored free magnetic energy that probably drives a CME can take many forms including (predominantly) mechanical in the form of an expanding CME and erupting filament, electromagnetic emission in the form of a flare, and also in the acceleration of energetic particles, magnetic field reconfiguration and bulk plasma motion. We mentioned the energy budget of CMEs and associated phenomena earlier: the few reports that have discussed this are Canfield et al. (1980), Webb et al. (1980), Emslie et al. (2004, 2005), and most recently, Ravindra and Howard (2010).
As noted in the introduction (Section 1), the formation and early development of CMEs can now be observed and followed near the surface, especially in EUV images (see, e.g., Hudson and Cliver, 2001). The first well-defined observation of the initial development of a CME “bubble” was with SOHO EIT images by (Dere et al., 1997). More recently, the high temporal and spatial resolution STEREO COR and EUVI and SDO AIA imagery has been used to study the initial formation of CMEs (see, e.g., Patsourakos et al., 2010a,b; Vourlidas et al., 2012). These demonstrate the rapid expansion of CMEs and their associated cavities and flux ropes.Update
EUV spectral observations from the UVCS, CDS, and SUMER instruments on SOHO and the SOT and EIS instruments on Hinode have helped us to measure the densities, temperatures, ionization states, and Doppler velocities of CMEs (e.g., Raymond, 2002; Kohl et al., 2006; Landi et al., 2010). Table 2 is a summary of the spectral lines that have been observed in CMEs by the UVCS instrument (Kohl et al., 2006). The UVCS instrument is unique in that it can sample the CME material at relatively high heights, e.g., out to , in the corona compared to the other spectrometers. Most CME material observed in UVCS is cool (< 105 K) and concentrated in small regions (Akmal et al., 2001), although this is not the case for fast CMEs associated with X-class flares (Raymond et al., 2003). Heating rates inferred from models using UVCS observations show that heating of the material continues out to and is comparable to the kinetic and gravitational potential energies gained by the CMEs (Akmal et al., 2001; Landi et al., 2009). The Doppler information from UVCS combined with the EIT and LASCO images has shown in one case the unwinding of a helical structure (Ciaravella et al., 2000). Doppler shifts are usually high, 1000 km–1, within halo CMEs, where compressed or deflected coronal material along the flanks of a CME is measured. H I Ly emission also suggests that dense material is present (Kohl et al., 2006).
|Line||Wavelength (nm)||log10 Tmax||Comments|
|H I Ly||121.567||4.3||radiative pumping|
|H I Ly||102.572||4.3||radiative pumping|
|H I Ly||97.254||4.3||radiative pumping|
|H I Ly||94.974||4.3||radiative pumping|
|C II||103.634, 103.702||4.3|
|N III||98.979, 99.158||4.8|
|N V||123.82, 124.280||5.3|
|O VI||103.191, 103.761||5.5||radiative pumping|
|Mg X||60.976, 62.493||6.1|
|Si III||120.651, 130.332||4.4||temperature-sensitive|
|Si XII||49.937, 52.066||6.3|
Living Rev. Solar Phys. 9, (2012), 3
This work is licensed under a Creative Commons License.