Several years later, MacQueen and Fisher (1983) analyzed the speed-height plots of 12 loop-like CMEs over the range observed by the MK3 coronagraph at the Mauna Loa Solar Observatory, and proposed the concept of two types of CMEs, i.e., flare-associated events that exhibit higher speeds (and show little acceleration with height) and filament eruption-associated events that exhibit slower speeds (and show large accelerations up to 50 m s–2). Sheeley Jr et al. (1999) used a new method to construct height-time maps of CMEs that were observed by SOHO/LASCO coronagraph. Similarly they classified CMEs into two types:
(1) Gradual CMEs: which are apparently formed when prominences and their cavities rise up from below coronal streamers, with speeds in the range 400 – 600 km s–1. They are often accelerating in the field of view of the coronagraphs;
(2) Impulsive CMEs: which are often associated with flares and Moreton waves on the visible disk, with speeds typically in excess of 750 km s–1. They are often decelerating in the field of view of the coronagraphs.
Although Cane et al. (1986) had already pointed out that quiescent filament eruptions are not necessarily associated with slow CMEs, the concept of two types of CMEs in association with flares or filaments was widely recognized (St Cyr et al., 1999; Andrews and Howard, 2001). For instance, Moon et al. (2002) performed a statistical study of 3217 CMEs observed by SOHO/LASCO during 1996 – 2000, and found that the fraction of flare-associated CMEs has a tendency to increase with the CME speed, whereas the fraction of eruptive-filament-associated CMEs tends to decrease with the CME speed. They claim that such a result supports the classification of the two types of CMEs, although the accelerations of the two types of CMEs are both concentrated near 0.
In order to explain the different kinematics, Low and Zhang (2002) phenomenologically proposed an idea for the two types of CMEs, i.e., the normal polarity flux rope eruptions correspond to the fast CMEs, whereas the inverse polarity flux rope eruptions correspond to the slow CMEs, as shown by Figure 37. On the other hand, it was found that, even for the inverse type only, the CME speed can be high or low (Chen and Krall, 2003; Wu et al., 2004). Whether the CMEs of the inverse polarity type would be statistically slower than the other type requires further observational and simulational support.
However, recently, Vršnak et al. (2005) presented a statistical analysis of 545 flare-associated CMEs and 104 non-flare CMEs in the heliocentric distance range of . They found that there is no distinct difference between filament-associated CMEs and flare-associated CMEs. A similar conclusion was drawn by Chen et al. (2005a). With a sample of 4315 CME events, Yurchyshyn et al. (2005) also found that there was no evidence to support the claim on two distinct types of CMEs. More and more evidence seems to indicate that the physics in the two types of CMEs is the same, with only the evolution time scale changing for individual events (Cliver and Hudson, 2002), i.e., for the faster CMEs, their characteristic time scales, e.g., the Alfvén transit time, are shorter, so the CMEs reach a high speed more rapidly.
The validation of the traditional classification of CMEs was further questioned by Feynman and Ruzmaikin (2004), who presented a CME associated with both a solar flare and an erupting filament. The analysis of this event and others led them to suggest that the apparent differences separating the two types of CMEs may be an observational effect, and all CMEs can be described by a single process as presented in Zhang et al. (2001a); see also Kahler et al. (1988) and Chen and Shibata (2000). Moreover, Chen et al. (2006a) pointed out that the traditional two-type classification, flare or filament associated, is logically incomplete since quite a lot of filament eruptions are accompanied by two-ribbon flares (e.g., Munro et al. 1979; see Zhang et al. 2001b for a typical event), which actually is the observational basis of the classical CSHKP model for CME/flares. Therefore, a complete classification should include the intermediate type of CMEs that are associated with both filament eruptions and flares. Therefore, they divide the events into three types in order to re-examine the CME velocity distributions: (1) CMEs associated with filament eruptions solely; (2) CMEs associated with flares solely; and (3) CMEs associated with both filament eruptions and solar flares. It was found that the average speeds of the three types of CMEs are 526 km s–1, 564 km s–1, and 738 km s–1, respectively. The Kolmogorov–Smirnov test indicates that the P-value for the likelihood between the velocity distributions of these three types is very high, suggesting that there is no significant difference between these types of CMEs.
Several factors contributed to reaching the conclusion of the existence of flare-associated and filament-associated CMEs by early researches. The first, as pointed out by Vršnak et al. (2005), is that the samples used by them were too small to be statistically reliable. The second factor is that some prominences are rooted behind the solar limb. For these events, the low-lying flares can not be observed. Another factor is that some CMEs are associated with SXR giant arcades (Hiei et al., 1993), which are physically the same as solar flares but can not be detected by GOES satellites (Shibata, 1996). Hence, the giant arcade-associated events are often categorized into the non-flare CMEs. For instance, the 1997 January 6 CME event, which was widely claimed to be not associated with any flare, was shown to correspond to a tiny flare with the GOES classification slightly below A-class (Wu et al., 2002).
Even though, CMEs do show diverse velocities, which still provokes an important question: what factors determine the CME velocity? By analyzing 13 events, Qiu and Yurchyshyn (2005) found that the CME velocity is proportional to the total reconnection flux, whereas Chen et al. (2006a) found that the CME velocity is better correlated to the average strength of the magnetic field in the filament channel. These results are roughly consistent with the classical CSHKP model, where magnetic reconnection is supposed to occur below a filament or flux rope. The Lorentz force of the reconnected field lines accelerates the reconnection outflow to the Alfvén speed at the inflow region. The upward reconnection outflow pushes the filament or flux rope to erupt. The erupting velocity of the flux rope might be proportional to the reconnection outflow speed. Since the Alfvén speed is proportional to the magnetic field strength, it is not surprising that the CME velocity is roughly linearly proportional to the magnetic field strength. While, in both papers, the strong correlation is based on small samples, where reconnection probably plays the main role in accelerating these CMEs.
If the CME dynamics is mainly determined by magnetic reconnection, we might see a remarkable correlation between CME velocity and the intensity of the flare, and the CME velocity is mainly determined by the magnetic field strength of the source region. However, Hundhausen (1997), Yashiro et al. (2002), and Vršnak et al. (2005) revealed that there is only a weak correlation between the CME apparent velocity and the peak flux of the associated flares. The situation becomes even worse if the CME velocities are corrected for the projection effects (Yeh et al., 2005). This implies that some factors, in addition to reconnection, may affect the CME velocity. One factor is the filament and its material drainage. As the evidence for this, if filament-associated CME events are excluded, the relation between the CME speed and the flare intensity is significantly increased (Chen and Zong, 2009). Another major factor might be the restraining force of the background field overlying the CME, as proposed by Török and Kliem (2007), who found that a steep decay of the background field with height leads to a fast CME and a gentle decay leads to a slow CME.
Living Rev. Solar Phys. 8, (2011), 1
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