1 Introduction

Solar flares are explosive phenomena observed in the solar atmosphere filled with magnetized plasma. The energy released during a flare is about 1028 – 1032 erg, and it takes various forms such as radiative energy, kinetic bulk energy, thermal and nonthermal energy. Because of their magnificent behavior, flares have been one of the most attractive scientific targets of the solar physics since they were first observed in the 19th century. The spatial size of a flare depends on individual events; in the smallest event the height of a flaring loop is less than 104 km, whereas it reaches 105 km in the largest event (see Section 2). The size also affects the duration of a flare (103 – 104 s) and the amount of energy released during a flare mentioned above.

Flares are observed in a wide range of electromagnetic waves such as radio, visible light, X-rays, and gamma rays (see Section 5). Emissions in these wavelengths come from the atmospheric layers extending from the chromosphere to the corona. In the extreme case, even the photosphere responds to a big flare, observed as white-light brightenings. Also a flare produces high-energy particles which travel through the interplanetary space, sometimes having a severe impact on the environment of the Earth (see Section 7).

Historically, flares were discovered in white light (Carrington, 1859; Hodgson, 1859). Later, spectroheliographs developed and an Hα filter was invented, flares can be observed in Hα. An Hα monochromatic image of a flare often shows beautiful two ribbons of bright patches, and the distance between these ribbons increases with time (e.g., Švestka, 1976Jump To The Next Citation Point; Zirin, 1988Jump To The Next Citation Point). For a long time flares were considered as chromospheric phenomena observed in Hα. However, the discovery of coronal radio and X-ray emissions from a flaring site has revealed that flares are actually coronal phenomena.

View Image

Figure 1: Two ribbons of a flare observed in Hα (I – V) and EUV (VI – X). Hα data are taken at Kwasan Observatory of Kyoto University, and EUV data are taken with EUV telescope aboard TRACE (from Asai et al., 2003).

Since the discovery of magnetic field on the Sun (Hale, 1908), the role of magnetic field in solar activity has been investigated extensively. Skylab mission (1973 – 1974) first made a detailed survey of the corona from the space using a soft X-ray telescope, revealing that bright regions in soft X-ray are well correlated with intense magnetic field. It is now widely accepted that the magnetic field provides a main energy source of the solar activity including flares. Observations suggest that the total magnetic energy stored in a typical sunspot with the size L and the averaged magnetic field B is

2 3 33( B )2( L )3 Emag ≃ (B ∕8π)L ≃ 10 103-G- 3-×-109-cm- erg, (1 )
which is sufficient to produce even the largest flare, although only a small portion of this total energy can be used, that is, a large amount of energy is unavailable because it is distributed as the potential field energy.

Following observational results, theoretical studies began to focus on the role of magnetic field in producing a flare. Giovanelli (1946) first pointed out that a neutral point where the magnetic field takes an X-type configuration could be the site of energy release during a flare. He proposed that the electric current may be dissipated at the neutral point. Hoyle (1949) also presented a similar idea about the mechanism for producing a flare, which is now known as magnetic reconnection. As Cowling (1953) later pointed out, however, there was difficulty in explaining the time scale of a flare if we assume that this is given by simple diffusion of magnetic field (no contribution of plasma flow). This is because the time scale of diffusion is much longer than the typical time scale of a flare in the corona where both the temperature of plasma (T) and the length scale of magnetic field (L) take a large value. The diffusion time in the corona is given by

2 14 9 2 6 3∕2 tdif ≃ L ∕η ≃ 10 (L ∕10 cm) (T ∕10 K) s, (2 )
4( --T---)−3∕2 2 −1 η ≃ 10 106 K cm s (3 )
is the magnetic diffusivity derived from the Spitzer resistivity (Spitzer, 1962Jump To The Next Citation Point). In order to overcome this big gap of these two time scales, the spatial scale of a diffusion region where electric current is dissipated must be as small as L ∼ 103 cm, although the total energy released during a flare cannot be explained by such a small diffusion region. Later, a model of magnetic reconnection that takes plasma flow into account was proposed by Sweet (1958Jump To The Next Citation Point) and Parker (1957Jump To The Next Citation Point). However, this reconnection model, which is now known as the Sweet–Parker reconnection, still cannot explain the time scale of a flare (Parker, 1963Jump To The Next Citation Point). This problem was partially solved by Petschek (1964Jump To The Next Citation Point)’s remarkable idea in which slow magnetohydrodynamic (MHD) standing shock waves are introduced to the reconnection dynamics (Petschek model, see Section 4). Since then, the magnetic reconnection has been considered to be one of the promising mechanisms for producing a flare, although a complete understanding of the relevant physics is still on the way.
View Image

Figure 2: Phenomenological models for flares based on magnetic reconnection. Lines with arrows on them indicate magnetic field lines. (a) Carmichael (1964Jump To The Next Citation Point), (b) Sturrock (1966Jump To The Next Citation Point), (c) Hirayama (1974Jump To The Next Citation Point), (d) Kopp and Pneuman (1976Jump To The Next Citation Point).

Several classic models based on magnetic reconnection have been proposed to explain the phenomenological aspect of flares: Carmichael (1964), Sturrock (1966), Hirayama (1974Jump To The Next Citation Point), and Kopp and Pneuman (1976Jump To The Next Citation Point) (see Figure 2View Image). These models assume more or less a similar configuration of magnetic field and its dynamic process, so these models are called with a single name, CSHKP model (Švestka and Cliver, 1992Jump To The Next Citation Point; Sturrock, 1992Jump To The Next Citation Point; Shibata, 1999Jump To The Next Citation Point). The CSHKP model has been a standard model of flares, and the basic features of this model are explained in Figure 3View Image. We also show a recent observational image obtained by Yohkoh in this figure.

Here we show a brief history about the naming of “CSHKP”. In the 1980s, this model was called the “Kopp–Pneuman model” in the United States. Shibata (1991) proposed to change the name from Kopp–Pneuman model to “SHKP model”, by respecting the pioneering works by Sturrock and Hirayama. Soon after, Sturrock himself proposed to add “C” in front of “SHKP” in his 1992 proceedings (Sturrock, 1992), noting another pioneering work by Carmichael. Švestka and Cliver (1992) also use the term “CSHKP model” in the same proceedings book.

The CSHKP model has been improved significantly since Kopp and Pneuman (1976Jump To The Next Citation Point), especially because the theory of magnetic reconnection has been developed, which is summarized in a nice review paper by Priest and Forbes (2002Jump To The Next Citation Point). Kopp and Pneuman (1976) considered that open field lines reconnect to form a closed loop, then the upflows of the solar wind along the reconnected field lines collide each other to generate a shock inside the closed loop, thereby heating coronal plasma up to the typical temperature of a flare. However, Cargill and Priest (1982Jump To The Next Citation Point) later pointed out that it is important to consider the slow MHD shocks suggested by Petschek (1964Jump To The Next Citation Point). Forbes and Priest (1984Jump To The Next Citation Point) noted that a fast MHD shock (termination shock) is also formed above the reconnected loop when a reconnection jet collides with this loop. Furthermore, Forbes and Malherbe (1986Jump To The Next Citation Point) showed that the slow MHD shock could be dissociated to an isothermal slow shock and a conduction front in the case of flares. Those features mentioned above have been reproduced by recent MHD simulations (Yokoyama and Shibata, 1997Jump To The Next Citation Point, 1998Jump To The Next Citation Point, 2001Jump To The Next Citation Point), which are discussed in Section 5.

As we mentioned, the theory of magnetic reconnection is not completed yet, not only in solar plasma but also in laboratory and magnetospheric plasma. This is why many anti-reconnection models of a flare have been proposed (e.g., Alfvén and Carlqvist, 1967Jump To The Next Citation Point; Akasofu, 1984Jump To The Next Citation Point; Melrose, 1997Jump To The Next Citation Point; Uchida and Shibata, 1988Jump To The Next Citation Point, etc.). However, recent space missions such as Yohkoh (1991 – 2001), SOHO (SOlar Heliospheric Observatory; 1995 –), TRACE (Transition Region and Coronal Explorer; 1998 –), RHESSI (Ramaty High Energy Solar Spectroscopic Imager; 2002 –), and Hinode (2006 –) have provided a huge amount of observational results suggesting that the magnetic reconnection occurs during a flare (e.g., Shibata, 1999Jump To The Next Citation Point). These observational results supporting the reconnection model of a flare are summarized in Section 3.

View Image

Figure 3: (a) Soft X-ray image of a long-durational-event (LDE) flare (see Section 2) observed by Yohkoh. (b) Schematic picture of a modified version of the CSHKP model, incorporating the new features discovered by Yohkoh (from Shibata et al., 1995Jump To The Next Citation Point).

The magnetic field providing the energy source of a flare originally comes below the solar surface (Zwaan, 1985Jump To The Next Citation Point), after traveling across the convection zone. While it travels across this region, the magnetic field is surrounded by a high-pressure plasma doing convective motions (e.g., Parker, 1979Jump To The Next Citation Point), so the magnetic field may take the form of a thin flux tube with some twist (Fan, 2009). When such a flux tube emerges into the surface, the background gas pressure decreases abruptly, so the magnetic field expand rapidly. The emerging magnetic field eventually fills a large volume up in the solar atmosphere and forms a magnetic structure there. The magnetic field loses the magnetic energy originally stored in the flux tube through this process, although part of magnetic energy remains in the magnetic structure as field-aligned electric current which does not produce the Lorentz force. The field-aligned current is therefore stored as free energy unless it is dissipated. The field-aligned current also introduces distortion into the magnetic structure (Sakurai, 1979Jump To The Next Citation Point), which is observed as a sheared arcade and/or twisted flux rope. The formation of distorted magnetic structure is important for understanding the state that leads to the onset of a flare, and we discuss this topic in Section 3.

It has been known that there are events preceding the onset of a flare. These are called ‘precursors’, and one of the prominent precursors is a newly emerging bipolar region at the surface, which may interact with preexisting magnetic field in the corona and produce a flare (Rust, 1968; Martres et al., 1968; Zirin, 1974; Heyvaerts et al., 1977Jump To The Next Citation Point; Martin et al., 1982; Feynman and Martin, 1995). Another well-known precursor is the activation, or eruption of a filament. A filament is composed of a relatively cool plasma with T ∼ 104 K, floated in the corona that is occupied by a much hotter plasma with 6 T ≥ 10 K. The activation of a filament suggests that some destabilization proceeds in a magnetic structure containing a filament. Since a filament is formed in a low-beta coronal region (magnetic pressure is dominant compared to gas pressure), the main forces causing such destabilization are the gradient force of magnetic pressure, magnetic tension force, and gravitational force. It is therefore important to know how these three forces keep balance during the pre-eruptive phase of a filament, and under what condition the balance is lost to cause the eruption of a filament. This issue is related to the onset of a flare and discussed in (Section 3).

Since the corona is filled with a highly conductive medium, dissipating electric current there is usually inefficient. This, on the other hand, indicates that the magnetic energy tends to be stored in the corona without any easy, continuous dissipation of electric current. An important question then arises regarding how to release the magnetic energy in the corona (otherwise tremendous amount of magnetic energy would be accumulated in the corona). One of the possible scenarios is that the magnetic energy is released via the dissipation of electric current in a thin, sheet-like region where the current density is enhanced, which is called current sheet. This is accompanied by the reconnection of magnetic field that converts the magnetic energy to not only thermal but also kinetic energy of plasma jets (high-energy particles are produced as well; selected particles are accelerated by electromagnetic force). The magnetic reconnection changes the configuration of magnetic field (field line topology) in a magnetic structure, destroying force balance to drive a global evolution of the structure. It has been suggested that a number of current sheets are spontaneously formed in the corona (Parker, 1994Jump To The Next Citation Point; Longcope, 2005Jump To The Next Citation Point; Low, 2006). Magnetic reconnection is a process of converting energies and changing topologies, so the coronal field can relax via successive reconnections to a lower energy state, avoiding a monotonous increase of the magnetic energy in the corona. We discuss the formation of a current sheet followed by rapid energy release by magnetic reconnection in Section 4.

The energy released by magnetic reconnection is immediately transported from the site of reconnection via radiation, thermal conduction, high-energy particles, and plasma blobs (Priest and Forbes, 2002Jump To The Next Citation Point). The configuration of magnetic field significantly influences this energy transport. Part of the released energy is transported downward along magnetic field lines via thermal conduction and high-energy particles, heating chromospheric plasma. This produces not only Hα ribbons shown in Figure 1View Image, but also hard X-ray (HXR) kernels in some energetic flares. The gas pressure of the heated chromospheric plasma increases, which drives the upflow of a plasma into the corona against gravitational force (chromospheric evaporation), filling a magnetic loop with hot plasma. This loop is observed in soft X-ray, so it is called a soft X-ray (SXR) loop. Another part of the released energy is transported upward as an ejecting plasma blob called plasmoid. During this phase selected particles are accelerated to become nonthermal energetic particles (Ramaty and Murphy, 1987Jump To The Next Citation Point). Those processes mentioned above are highly dynamic and complex, so numerical simulations are quite useful and indispensable for a better modeling of time-dependent, nonlinear energy transporting processes. In Section 5, we show the result obtained from such numerical simulations and explain key issues of energy transport during a flare.

A flare is often associated with other dynamic phenomena on the Sun. As we mentioned, the activation of a filament is one of those flare-associated phenomena. Sometimes a flare is part of a large-scale eruption known as a coronal mass ejection (CME), which carries a tremendous amount of plasma (up to 1016 g) to the interplanetary space. In a typical CME, a large magnetic loop with the size of the solar radius moves away from the Sun at 30 – 2500 km s–1 (Yashiro et al., 2004), which forms a shock wave at the front of the erupting loop. Regarding the modeling of CME, there are several concise and comprehensive reviews (Forbes, 2000Jump To The Next Citation Point; Klimchuk, 2001; Gibson et al., 2006; Mikić and Lee, 2006; Forbes et al., 2006; Chen, 2011). The physical relationship between a flare and CME has intensively been investigated (Gosling, 1997).

It should be noted that this paper is mainly focused on the MHD aspects of a flare. Hence, for readers who are interested in other aspects of a flare such as nonthermal processes of particles we introduce the following papers and books: Švestka (1976Jump To The Next Citation Point), Priest (1981), Dulk (1985), Dennis (1985Jump To The Next Citation Point), Sakai and Ohsawa (1987Jump To The Next Citation Point), Tandberg-Hanssen and Emslie (1988), Haisch et al. (1991), Kahler (1992), Hudson and Ryan (1995), Sakai and de Jager (1996), Low (1996Jump To The Next Citation Point), Low (2001), Shibata (1997), Shibata (1999Jump To The Next Citation Point), Miller (1997), Bastian et al. (1998), Hundhausen (1999), Forbes (2000), Charbonneau et al. (2001Jump To The Next Citation Point), Aschwanden (2002Jump To The Next Citation Point), Priest and Forbes (2002), Zhang and Low (2003), Lin et al. (2003). The readers can also find a well-described explanation of solar flares in the following books: Parker (1979), Parker (1994), Melrose (1980), Priest (1982), Zirin (1988), Stix (2004), Benz (1993), Tajima and Shibata (1997Jump To The Next Citation Point), Golub and Pasachoff (1997), Somov (2000), Priest and Forbes (2000Jump To The Next Citation Point), Aschwanden (2004Jump To The Next Citation Point). Let us show the organization of this paper. In Section 2, we summarize observational features of a flare and flare-like phenomena, and show their phenomenological models. Section 3 describes energy build-up via flux emergence which plays the central role in forming a sheared magnetic structure with the free energy stored inside it. In this section we discuss the dynamics of an emerging flux tube to derive the basic feature of a magnetic structure evolving into the onset of a flare such as a filament and sigmoid. In this respect, force-free modeling and magnetic helicity are also discussed. Section 4 is focused on magnetic reconnection, the central engine that converts free energy to produce various dynamic processes. We start with the basic concept of magnetic reconnection and then demonstrate its application in the physics of flares. In Section 5, we show various processes associated with energy transport in a flare, such as radiation, mass ejection, shock heating, wave propagation, and particle acceleration. In Section 6, stellar flares are briefly explained. Finally, we make some concluding remarks on the physical processes responsible for producing a flare. The contents of the paper are listed below.




Section 2

Observational features and phenomenological models

LDE flares, Impulsive flares, Giant arcade, Transient brightenings, X-ray jets, and models focused on their phenomenological aspects

Section 3

Energy build-up

Flux emergence, Force-free fields, Preflare structure, Magnetic helicity

Section 4

Energy release

Magnetic reconnection, Current-sheet formation, Flux-rope eruption

Section 5

Energy transport

Radiation, Mass ejection, Shock heating, Particle acceleration, Wave propagation, Chromospheric evaporation

Section 6

Stellar flares

Application of solar flare model to stellar flares

Section 7


Summary of physical processes for producing a flare

  Go to previous page Go up Go to next page