Solar flares
Solar flare soft X-ray burst classification
The current standard measure for solar flare intensity is the widely-used Geostationary Environmental Satellite System (GOES) ABCMX SXR classification system which is defined as follows: SXR classes A1-9 through X1-9 correspond to flare peak 1–8 Å fluxes of 1–9 × 10−n W m−2 where n = 8, 7, 6, 5, and 4, for classes A, B, C, M, and X, respectively. Occasionally, flares are observed which peak intensities ≥ 10−3 W m−2. Rather than being assigned a separate letter designation, such flares are referred to as X10 events and above.Footnote 4
Approximately 20 flares of class X10 or higher have been observed during the last ~ 40 years.Footnote 5 The largest GOES SXR flare yet recorded occurred on 4 November 2003. The 1–8 Å detector saturated at a level of X18.4, with an estimated SXR class of X35 ± 5 (Cliver and Dietrich 2013) based on consideration of values given in Kiplinger and Garcia (2004), Thomson et al. (2004, 2005), Brodrick et al. (2005), and Tranquille et al. (2009). The bulk of the SXR class estimates for this flare were obtained via comparisons of flare SXR intensities and the amplitudes of flare-associated sudden ionospheric disturbances (SIDs; Mitra 1974; Prölls 2004; Tsurutani et al. 2009) caused by energetic flare photons leading to increased electrical conductivity in the day-side ionosphere, e.g., the magnetic crochet observed for the 1859 event. Cliver and Dietrich (2013) obtained an estimate for the Carrington flare of X45 ± 5 based on the magnetic crochet (a type of SID recorded by ground-based magnetometers) observed for this event (Cliver and Svalgaard 2004; Boteler 2006; Clarke et al. 2010). The November 2003 and September 1859 flares provide the current benchmarks for extreme flare activity. Less certain is an estimate of X285 ± 140 (Cliver et al. 2020b: Sect. 7.8 below) for the flare associated with the inferred SEP event of 774 AD (Miyake et al. 2012; Usoskin et al. 2013). Because flares such as those observed/inferred in 1859 and 774 AD are rare, we need to look at the ensemble of lesser flares observed by the GOES system since 1976 to estimate the occurrence frequency of extreme flares of this size and larger.
Solar flare frequency distributions
The GOES 1−8 Å soft X-ray measurements provide the longest, uninterrupted, and uniformly calibrated data set available for solar flares. Thus these records provide the basis for statistical studies that seek to establish the frequency of occurrence of flares as a function of flare magnitude. The most often used magnitude metric is the peak intensity of the flare as measured in that passband (expressed in W m−2), i.e., the ABCMX GOES classification.
Flares of classes A and B are under-reported in the GOES records because during active phases of the solar cycle such faint flares are often hard to separate from (or detect above) the supposedly quiescent background and are judged to be of low interest. As a further complication, flare magnitude is often not corrected for background emission, so that the weak flares that are reported are intrinsically overestimated in their strength. The least observationally biased records are those for flares of classes C, M, and X.
Many studies over the past decades have established that, at least for the larger flares, the frequency distribution for peak brightness of flares is well approximated by a power law (a thorough review in the literature of power-law fits to flare strength parameters can be found in Aschwanden et al. (2016) who also discuss waiting time distributions). The first such distribution listed by Aschwanden is that of Akabane (1956) who presented a distribution for burst peak radio intensities at 3.0 GHz. Subsequently, Hudson et al. (1969) showed a frequency histogram based on solar hard X-rays as observed by the Orbiting Solar Observatory (OSO-3) satellite that, with hindsight, hints at a power-law distribution for the more intense flares, although it lacks an explicit power-law fit to the data. Such an approximation was made by Drake (1971) based on soft X-ray measurements by the Explorer 33 and 35 spacecraft. Similar power-law distributions were first reported a few years earlier for UV Ceti type flare stars (very cool main sequence stars that exhibit frequent flaring; e.g., Kunkel 1968; Gershberg 1972; Lacy et al. 1976; Kowalski et al. 2013).
Rosner and Vaiana (1978) suggested that one way to create such frequency distributions would be to have a system in which exponential growth in stored energy is interrupted at random times by a flare-like event in which a substantial amount of the stored energy is removed from the system. But if such an energy build-up would occur in regions on the solar surface, a correlation between flare brightness and waiting time would be expected. The absence of such a correlation in both solar observations (Aschwanden et al. 1998) and stellar data (as already noted by, e.g., Lacy et al. 1976) led to other ideas, including that of self-organized criticality (SOC; Bak et al. 1987, 1988; applied to solar flares by Lu and Hamilton 1991). The SOC model for solar flares is based on the assumption that flares occur as a time-varying, or non-stationary, Poisson process. We point the reader to Aschwanden et al. (1998, 2016) for descriptions of the developments of such interpretations, focusing here on the distribution functions rather than the processes behind them.
In an extensive study of ~ 50,000 GOES SXR bursts observed over a 25-year period from 1976 through 2000–38,000 of which were in classes C-X, Veronig et al. (2002a) obtained well-defined fluence and peak flux distributions that could be fitted by power laws with similar exponents: − 2.03 ± 0.09 and − 2.11 ± 0.13, respectively. These power-law approximations to the frequency distribution hold for over 2.5 orders of magnitude, with no significant indication of a change in behavior up to the largest flares observed. At a flare magnitude bin at ~ X15 (based on only two flares), the histogram simply aborts at the end of a straight power law (Fig. 7).
The work by Gopalswamy (2018) includes a comparable analysis of flare magnitudes over a 47-year period from 1969 through 2016, including 55,285 flares of class C1 or larger (it did not consider the ~ 11,500 B class flares included in the Veronig et al. 2002a, sample). This larger sample was obtained over a time base that almost doubles that of Veronig et al. (2002a) by including more recent data and also extending the GOES records with Solar Radiation (SOLRAD) satellite data for the period 1969–1975. The downward-cumulative representation of the frequency distribution for that data set (Fig. 8) exhibits a deviation from a power-law distribution (see also, e.g., Riley 2012). In a log–log downward cumulative representation, an approximation by power laws suggests a break towards less-frequent larger flares somewhere around X4, defined by a total of about 100 flares at that magnitude or larger. Such a downward break is consistent with the conclusions reached by Schrijver et al. (2012).
The century- and millennium-level flares based on the modified exponential function fit to the SXR peak values in Fig. 8 are X44 and X101, respectively. The corresponding bolometric energy values from Gopalswamy (2018) are ~ 4 × 1032 erg and ~ 1033 erg. The 100- and 1000-year estimates based on power laws are comparable (X42 and X115; ~ 4 × 1032 erg and ~ 1.2 × 1033 erg). Both functional forms fit the tail of the distribution well. For a 10,000-year flare, the exponential function yields ~ X200 versus ~ X310 for the power-law fit.
While the radiative energy of the strongest observed solar flare is often given variously as 1032 erg, about 1032 erg, or ~ 1032 erg (e.g., Candelaresi et al. 2014; Peter et al. 2014; Osten and Wolk 2015; Hudson 2016; Maehara et al. 2017; Notsu et al. 2019; Brasseur et al. 2019), the measured value is a half-decade larger. Emslie et al. (2012) reported that the inferred/observed radiative energies of 12 events from 2002 to 2006 were above the 1032 erg threshold. For the three largest such events (4 November 2003, 4.3 × 1032 erg, X35 ± 5; 28 October 2003, 3.6 × 1032 erg, X17; 7 September 2005, 3.2 × 1032 erg, X17), the bolometric energy was measured directly by the Total Irradiance Monitor (TIM; Kopp and Lawrence 2005) on the Solar Radiation and Climate Experiment (SORCE; Rottman 2005). Thus, the largest observed (November 2003) and inferred (September 1859; ~ 5 × 1032 erg, ensemble estimate of X45 ± 5; Cliver and Dietrich 2013; based on electromagnetic emissions) flares are comparable to the 100-year flare given by Fig. 8.
White-light flares
The flare Carrington and Hodgson independently recorded on 1 September 1859 was observed in integrated light. Thus their Monthly Notices papers were the first reports of what came to be known as “white light flares” (WLFs). A compilation of reports of such events from (1859–1982) is given in Neidig and Cliver (1983a). At that time, Neidig and Cliver (1983b) reckoned that such flares were relatively rare, with an occurrence frequency of ~ 15 per year based on a ~ 2.5-year period from June 1980 to December 1982 following the maximum of solar cycle 21. They determined that a flare with SXR class ≥ X2 originating in a large (≥ 500 μsh), magnetically complex, sunspot group was a sufficient condition for a white light flare. Two decades later, Hudson et al. (2006) used Transition Region and Coronal Explorer (TRACE; Handy et al. 1999) images, which encompassed a wavelength range from 1500 to 4000 Å, and identified 11 white-light flares that had SXR classes ranging from C1.6 to M9.1 (median = M1.2) to “support the conclusion of Neidig (1989) that white-light continuum occurs in essentially all flares.” Subsequently—based on observations of many flares with the Variability of Solar Irradiance and Gravity Oscillations (VIRGO: Fröhlich et al. 1995) experiment on the Solar and Heliospheric Observatory (SOHO; Domingo et al. 1995)—Kretzschmar (2011) estimated that white-light emission might typically account for ~ 70% of the total radiated flare energy, following an earlier estimate by Neidig (1989) of ≲ 90% based on analysis of the 24 April 1981 white-light flare (Neidig 1983).
White-light flares are of particular importance in studies of extreme flares on Sun-like stars because of the observations of such flares by the high-precision (better than 0.01% for moderately bright stars) visible wavelength (i.e., white-light, 4000–9000 Å) photometer on the Kepler spacecraft (Koch et al. 2010), which operated from 2009 to 2018. The durations of stellar white-light flares can exceed several hours (Maehara et al. 2012) versus typical durations of ≲ 10 min (Neidig and Cliver 1983a) for their largest solar counterparts, although some of this discrepancy may be due to the lower sensitivity of early WLF observations (see Fig. 13 below). Stellar flares detected by Kepler have bolometric energies exceeding 1033 erg (Maehara et al. 2012; Shibayama et al. 2013), exceeding those of the largest solar flares (~ 4 × 1032 erg; Emslie et al. 2012; cf. Kane et al. 2005) by about a factor of three. Whereas the Kepler photometer detects essentially all superflares with amplitudes > 1% of the average stellar flux on solar-type (G-type dwarfs, 5100 K < Teff < 6000 K and log g > 4.0) stars, corresponding to a flare bolometric energy of ~ 5 × 1034 erg (Shibayama et al. 2013), the estimated detection completeness for ~ 1034 and 1033 erg flares are ~ 0.1 and 0.001, respectively (Maehara et al. 2012).
It is well-accepted that the paradigm of the solar flare also holds also for the large (1034–35 erg) stellar flares (Gershberg 2005). In the standard picture of solar flares, a rapid conversion occurs, via magnetic reconnection, of energy stored in the magnetic field to energies distributed over bulk kinetic energy and thermal and non-thermal particle distributions. A significant fraction of this converted energy is eventually released as a pulse in the visible radiation, the observable diagnostic of the Kepler superflares on Sun-like stars. Like solar flares, stellar flares: (1) occur on single “solar-type” stars, viz., G-type main sequence stars with Teff and gravity similar to that of the Sun and rotation periods > 10 days and without hot Jupiters (Maehara et al. 2012; Shibayama et al. 2013; Okamoto et al. 2021; Sect. 3.2.2); (2) show evidence of non-thermal particle populations and of temperatures ≳ 107 K (e.g., Osten et al. 2007; Benz and Güdel 2010); (3) exhibit a characteristic fast rise and slower, near-exponential, decay in X-rays (Kahler et al. 1982; Haisch et al. 1983); and (4) have lower-energy emissions that often scale with the time integral of their high-energy emissions (equivalent to the Neupert effect for solar flares; Hawley et al. 1995; Güdel 2002; Osten et al. 2004; see Sect. 3.1.4). The above evidence for the commonality of stellar flare characteristics and processes with those that occur during smaller solar flares implies that the statistics of energetic flares on a multitude of Sun-like stars (see Sect. 3.2) can provide information on extreme solar flares that may occur only once per millennium or even less frequently (Notsu et al. 2019; Okamoto et al. 2021).
Impulsive and gradual phases of flares: The Neupert effect
The separation of flare emission into a fast-rise impulsive phase followed by a slowly decaying gradual phase has long been noted in multiple wavelengths (Fletcher et al. 2011). The impulsive phase of flares (Dennis and Schwartz 1989; Benz 2017), characterized by non-thermal hard X-ray and radio bursts, marks the time of the early principal energy release in a flare. See Fletcher et al. (2011) and Benz (2017) for detailed reviews of flare observations.
In a prescient six-page paper based on early flare soft X-ray data for three flares, Neupert (1968) found that the integral of impulsive phase microwave emission in each flare resembled the rise of the SXR light curve. He hypothesized that collisional losses by the energetic electrons responsible for the microwave burst heated the chromosphere to sufficient temperatures to eject plasma into the low corona where its cooling was manifested by the slow decay of thermal SXR emission during the flare gradual phase.
The support structure for Neupert’s conjecture on the conversion of non-thermal energy to thermal energy in a solar flare lay in the future. It consisted of: the thick target model of hard X-ray emission (Brown 1971; Kane and Donnelly 1971), the CSHKP model for eruptive solar flares (Carmichael 1964; Sturrock 1968; Hirayama 1974; Kopp and Pneuman 1976; Hudson 2021), and the establishment of chromospheric evaporation (Antonucci et al. 1984; Fisher et al. 1985). In addition, evidence that white light emission is powered by energetic electrons accelerated during the flare impulsive phase accumulated (Hudson 1972; Machado and Rust 1974; Rust and Hegwer 1975; Neidig 1989; Hudson et al. 1992; Neidig and Kane, 1993). Figure 9 (taken from Hayes et al. 2016) shows the inverse aspect of the Neupert effect—hard X-ray emission is the derivative of the SXR time profile (Dennis and Zarro 1993)—using modern data.
Figure 10 shows a CSHKP schematic for an eruptive flare that illustrates various elements of the Neupert effect: acceleration of electrons and protons at a neutral current sheet, propagation of electrons to the low atmosphere giving rise to hard X-ray and microwave emission and heating of the chromosphere via particle bombardment, evaporation of the heated plasma to fill SXR emitting loops, and subsequent retraction and cooling of these loops as manifested by the appearance of post-flare Hα loops (Švestka et al. 1987).
According to the Neupert effect, the impulsive phase of flares can be interpreted as the time interval during which magnetic reconnection occurs. In the two-dimensional schematic in Fig. 10, this phase would end when the last closed loop of the CME is disconnected from the pinched off lower loop system. In theory, the SXR intensity at this time would be at it or near its maximum. In reality, Veronig et al. (2002b) found that this coincidence or near-coincidence of the end of the impulsive phase with SXR maximum was observed for only about 50% of their sample of ~ 1100 C2-X4 flares. In ~ 25% of the cases, the SXR emission peak occurred after the impulsive phase hard X-ray emission, and another ~ 25% of cases were unclear. For a much smaller subset of 66 large impulsive hard X-ray flares, Dennis and Zarro (1993) found that 80% of the events were consistent with the Neupert effect, suggesting that a relationship between SXR emission and white light emission can be valid for extreme events.
The significance of the Neupert effect for this review is two-fold: (1) it provides a link between the widely used (and more readily available) SXR intensity as a measure of flare size and the physically more significant measure of flare white-light energy which dominates the radiative energy budget of flares; and (2) as will be seen in Sect. 3.2, it allows flares on Sun-like stars to be directly compared with solar flares in terms of the standard CMX SXR classification.
The classic Neupert effect suggests that magnetic reconnection, particle acceleration, and post-flare loop formation, are confined to the flare impulsive phase, a useful generalization. Subsequently, Zhang et al. (2001) demonstrated that the SXR rise phase corresponded to the interval of CME acceleration. That said, magnetic reconnection, particle acceleration, and loop formation can continue for up to ≳ 10 h (e.g., Bruzek 1964; Kahler 1977; Akimov et al. 1996; Gallagher et al. 2002; Tripathi et al. 2004), well beyond the impulsive phase, as can particle acceleration by CME-driven shocks. Moreover, these late phase phenomena can give rise to strong, and even extreme (e.g., Frost and Dennis 1971; Chupp et al. 1987; Gary 2008; Hudson 2018) solar phenomena. Time profiles of flare emission at various frequencies/energies are given in Fig. 11. The hard and soft X-ray traces during the impulsive phase in panels (b) and (c) display the classic Neupert effect, which applies for the majority of flares, particularly the smaller confined events, but also for eruptive events with late phase reconnection too weak (or too high in the corona) to significantly affect the SXR time profile in panel (c)—consistent with the radiative energy domination of the impulsive phase. As shown in panels (b), (d), and (e), and discussed in Sect. 4, electrons accelerated in certain of these eruptive events can result in gradual hard X-ray and microwave bursts as well as intense decimetric bursts at ~ 1 GHz, with peak flux densities as large as 105–106 solar flux units.
Late phase particle acceleration in eruptive flares can also manifest itself in remotely-sensed high-energy solar γ-ray emission. When first observed in the 1980s and early 1990s, prolonged 100 MeV γ-ray emission, attributed to the decay of neutral pions which require acceleration of protons to ≳ 300 MeV energies—the highest that can be inferred from γ-ray observations—for their production, contrasted sharply with the Neupert effect in which flare energy degrades over time from non-thermal X-ray and radio emissions to thermal soft X-rays. Rather than flare electromagnetic emission becoming progressively less energetic after the impulsive phase, 100 MeV γ-ray emission characteristically occurs during the late phase of flares (Share et al. 2018). The most economical (Occam’s razor) explanation for acceleration of the energetic protons responsible for such emission is that they are accelerated by the same coronal shocks (manifested by the slow-drift metric type II burst in panel (a)) that give rise to the solar energetic particles detected by spacecraft near Earth (Sect. 7).
Energetics of flares and CMEs
1−8 Å soft X-ray and broad-band visible light emissions are but two of the channels into which the processes related to solar flares deposit energy. Magnetic energy is converted not only into photon emission in other channels, but also into kinetic energy of non-thermal particle populations, and the bulk kinetic energy associated with the motion of material ejected from the flaring active region. Specific electromagnetic, plasma, and particle aspects of solar eruptions will be considered in following sections of this review.
Figure 12 (taken from Emslie et al. 2012) shows the various pathways by which energy is converted and released during eruptive solar flares. Emslie et al. analyzed wide-ranging data for a sample of 38 M- and, mostly, X-class flares. Figure 12 summarizes their findings for six well-studied X-class flares. All of the listed phenomena in the figure derive their energy ultimately from the free (or excess) magnetic energy of an active region (Efm), defined to be the non-potential energy beyond (or in excess of) that obtained from a potential field model. Emslie et al. (2012) estimated that, on average, the amount of Efm of an active region was equal to 30% of the modelled potential energy determined from line-of-sight magnetograms. For the six events upon which Fig. 12 is based, with an average flare magnitude of X6, the median Efm was ~ 15 × 1032 erg. Emslie et al. estimated that ~ 30% of the free energy in an active region was released in an eruptive event. This, in turn, implies that the energy released in large flares amounts to ~ 10% of the magnetic energy in an active region (see also Shibata et al. 2013). Thus, it is not surprising that magnetic field changes between before and after a flare are hard to spot unambiguously in the overall, always-evolving, multi-thermal coronal configuration. In contrast, rapid changes associated with flares are detected in the photospheric field, most readily in association with SXR flares of class M or X (e.g., Wang et al. 2002; Liu et al. 2005; Sudol and Harvey 2005; Castellanos Durán et al. 2018). The search for such changes was lengthy, however, beginning with Carrington (1859) who looked unsuccessfully following the 1 September 1859 flare (Fig. 1) for any changes in the sunspot group based on the sketch he had made prior to the flare.
In keeping with the increased emphasis on CMEs versus flares per se for space weather during the last ~ 25 years, Emslie et al. (2012) find that, on average—for the two dominant terms of energy release (Fig. 12)—CME energy (kinetic plus gravitational potential) is ~ 3 times larger than the flare bolometric energy. Aschwanden et al. (2017) performed an analogous analysis to that of Emslie et al. (2012) for a sample of 157 M- and X-class eruptive flares observed from 2002–2006 and found that the ratio of CME to bolometric energy was ~ 1:1 versus ~ 3:1 determined by Emslie et al. (2012). These different samples preclude a straightforward direct comparison. X-class flares constituted 76% of the Emslie et al. (2012) sample with a median of X2 versus < 10% of that for Aschwanden et al. (2017; median < M2). As this review focuses on the largest, most energetic events, we use the Emslie et al. result.
This ~ 3:1 ratio CME to flare energy obtained by Emslie et al. is presumably an underestimate because the magnetic energy of the CME (Webb et al. 1980) is not taken into account in the determination. In addition, the CME masses and speeds in the SOHO Large Angle Spectroscopic Coronagraph (LASCO; Brueckner et al. 1995) CME data base (https://cdaw.gsfc.nasa.gov/CME_list/; Yashiro et al. 2004; Gopalswamy et al. 2009) used by Emslie et al. are underestimated because of projection effects (Vourlidas et al. 2002, 2010; Vršnak et al. 2007; Paouris et al. 2021). Emslie et al. (2012) write that the mass is underestimated by about a factor of two for CMEs that are associated with flares ≲ 40° from the limb—approximately half of their sample—with greater underestimates for those closer to disk center. Vršnak et al. find that average velocities of non-halo limb-CMEs are 1.5–2 times higher than for such CMEs originating near disk center. This underestimation of CME parameters impacts estimates of the largest possible solar flare (Sect. 3.1.7) that are based on active region energy and flare energy released by reducing the radiative energy budget. For a plausible 9:1 apportionment (vs. 3:1), only 10% of the released energy would go into the flare.
In reference to extreme solar flares, we can safely conclude that they are eruptive. Schrijver (2009) reviewed studies on flares with and without CMEs, including Andrews (2003), Yashiro et al. (2005), and Wang and Zhang (2007), that showed whereas ~ 20% of low to mid C-class flares have associated CMEs, ~ 90% of X-class flares are eruptive. The largest reported flares that lacked CMEs were an X3.1 flare on 24 October 2014 from NOAA spot group 12192 that produced several confined X-class flares (e.g., Thalmann et al. 2015; Liu et al. 2016; Green et al. 2018; Gopalswamy 2018), an X3.6 flare on 16 July 2004 (Wang and Zhang 2007), and an X4.0 flare on 9 March 1989 (Feynman and Hundhausen (1994). For flares on active stars, it has been argued (Drake et al. 2016; Alvarado-Gómez et al. 2018; Moschou et al. 2019; Li et al. 2020, 2021) that much stronger magnetic fields than observed on the Sun would not allow free energy to be released via eruption below some threshold. Based on observations of solar CMEs, Li et al. (2021) speculate that for an active region with unsigned flux of 1024 Mx on a solar- type star, the CME association rate for X100 flares would be < 50%. To date, the most intense solar flares yet inferred (viz., 1 September 1859 and 774 AD (Sect. 7.8)) are eruptive by virtue of their terrestrial effects—geomagnetic storm and proton event, respectively—which both imply CME association (Kahler 1992; Gosling 1993).
Empirical relationship between flare total solar irradiance and SXR class
Even though most of the solar flare electromagnetic radiation, or bolometric output, occurs in the visible range, the contrast with the quiescent photosphere is low. Only the larger white-light flares stand out clearly against the photospheric background in high-resolution images while only the most extreme flares such as 28 October or 4 November 2003 can be recognized as spikes in the record of total solar irradiance (TSI; Woods et al. 2004, 2006; Kretzschmar 2011). As shown in Fig. 13 (from Kopp 2016), the X17 flare on 28 October in the Halloween sequence in 2003 reached a peak brightness in the visible wavelength range of only 0.028% of TSI.
However, when averaging observations over multiple flares, the signal-to-noise ratio increases. Kretzschmar et al. (2010) and Kretzschmar (2011) used this ensemble technique in a superposed-epoch analysis of observations made with the VIRGO experiment on SOHO to correlate the GOES measured soft X-ray energy (\(\mathcal{F}\) GOES) and bolometric energy of over 2100 solar flares (with GOES classes from C4 up to X17) over an 11-year period. Kretzschmar (2011) presented his results averaged over fairly wide energy ranges in order to have sufficient signal-to-noise. In their analysis of solar and stellar flares, Schrijver et al. (2012) fitted a power-law relationship to these SOHO and GOES measurements to obtain
$$ {\mathcal{F}}_{{{\text{TSI}}}} = 2.4 \times 10^{12} {\mathcal{F}}_{{{\text{GOES}}}}^{0.65 \pm 0.05} , $$
(3)
for the conversion from 1 to 8 Å GOES radiated energy to bolometric energy (both expressed in erg) based on flares with bolometric energies ranging from 3.6 × 1030 erg to 5.9 × 1031 erg. The GOES 1−8 Å channel, commonly used to characterize flares and their frequency distribution, captures < 2% of the total radiation emitted by large flares (Woods et al. 2006; Kretzschmar 2011; Emslie et al. 2012).
To go from soft X-ray radiated energy to peak brightness (which sets the GOES flare class) we can use observations of the almost 50,000 flares that were analyzed by Veronig et al. (2002a). They find that the GOES 1−8 Å radiated energy,
$$ {\mathcal{F}}_{{{\text{GOES}}}} = 2.8 \times 10^{29} { }\left( {\frac{{{\mathcal{C}}_{{{\text{GOES}}}} }}{{{\mathcal{C}}_{{{\text{GOES}},{\text{ X}}1}} }}} \right)^{1.10} { ,} $$
(4)
with their SXR fluence measurement converted to energy by multiplying by 4π × (1 AU)2 and peak SXR flux scaled to equal unity for an X1 class flare (10−4 W m−2), where 1 AU = 1 astronomical unit (1.5 × 108 km).
Combining Eqs. (3) and (4) leads to a relationship between total radiated energy, \({\mathcal{F}}_{{{\text{TSI}}}}\) (in erg), and the scaled GOES flare class:
$$ {\mathcal{F}}_{{{\text{TSI}}}} = 0.33 \times 10^{32} { }\left( {\frac{{{\mathcal{C}}_{{{\text{GOES}}}} }}{{{\mathcal{C}}_{{{\text{GOES}},{\text{ X}}1}} }}} \right)^{0.72} . $$
(5)
Thus values of 1032 (1033) erg correspond to a ~ X5 (~ X115) flare, with uncertainties in the mean relationships as well as flare-to-flare differences resulting in overall uncertainties of easily half an order of magnitude (e.g., Emslie et al. 2012; Benz 2017; Gopalswamy 2018). Because TSI is dominated by white-light emission in the impulsive phase of flares, Eq. (5) can be considered a corollary of the Neupert effect.
Estimates of the largest possible solar flare based on the largest observed sunspot group
(a) Estimates based on reconnection flux
Large flares require large active regions with substantial amounts of magnetic flux near a high-gradient polarity separation line (Schrijver 2007). As noted in Sect. 2.1, the largest sunspot group recorded since systematic area measurements began at RGO in 1874 (Greenwich group 14886) occurred on 8 April 1947 (Fig. 14), with a corrected total spot area of 6132 μsh. Estimates of the total unsigned magnetic flux of such an active region range from 2 to 6 × 1023 Mx (Toriumi et al. 2017; Schrijver et al. 2012). Figure 15 (adapted from Toriumi et al.) shows a plot of unsigned magnetic flux (from spots and plage) for active regions associated with 51 flares located within 45° of solar central meridian from 2011–2015 (black diamonds). The red circle data points with fitted dashed line are based on unsigned flux values from Kazachenko et al. (2017). Linear fits to both data sets give a value of ~ 2 × 1023 Mx for the April 1947 spot group. At the same time, the largest spot groups in this figure (all but one from spot group 12192 in October 2014) hint at a non-linear fit line in the log–log plot that could reach the 6 × 1023 Mx value used by Tschernitz et al. The large black circle in Fig. 15 represents a compromise average flux value of 4 × 1023 Mx.
The Schrijver et al. estimate of ~ 6 × 1023 Mx assumes a uniform field strength of ~ 3000 G for the spotted area of group. While ~ 3000 G is not unreasonable for the peak field strength of a large single spot, such a value is reached < 10% of the time in sunspot umbrae (Pevtsov et al. 2014). Moreover, the corrected umbral area of Greenwich 14886, is only 739 μsh. Schrijver et al. (2012) used an overall sunspot group area of 6000 μsh but did not consider plage. This neglect of the plage contribution offsets the assumption of the high uniform field strength in the sunspots.
Cycle18 (1944–1954), which included the great spot group of April 1947 is known as the cycle of “giant” sunspots (Dodson et al. 1974). The five spot groups since 1874 with observed areas > 4500 μsh all occurred in this cycle. (The next largest group, with an area of 3716 μsh, occurred in January 1926.) While large spot areas are indicative of the potential for extreme events, they are not sufficient for their occurrence. Of these five spot groups, two (February 1946 (− 220 nT; http://dcx.oulu.fi/dldatadefinite.php) and July 1946 (− 268 nT) were associated with “great”, but not “outstanding” magnetic storms, the designation given the May 1921 event (Jones 1955). As we shall see in Sects. 6 and 7, the intensities of geomagnetic storms and solar energetic proton events, respectively, can be affected by other factors than the size of the parent sunspot group.
Tschernitz et al. (2018) find a strong correlation between the GOES class of a flare and the rate of magnetic reconnection for a sample of 51 flares (SXR class B3-X17). They estimate the reconnection rate by mapping the Hα ribbon evolution in time onto the region’s magnetogram (Fig. 16a). Among the most pronounced correlations (r = 0.92) they find a scaling between GOES SXR class and the total reconnected flux \({\Phi }_{\mathrm{r}}\) (in Mx; 1 Mx = 10−8 Wb) (Fig. 16b). Tschernitz et al. (2018) find that the largest events involve about half of the magnetic flux contained in the active region. For the April 1947 spot group with a total magnetic flux of 6 × 1023 Mx (after Schrijver et al. 2012), Tschernitz et al. inferred a largest possible SXR flare of class ~ X500 (with confidence bounds from X200-X1000) as shown in Fig. 16b, with a corresponding bolometric energy of ~ 3 × 1033 erg. However, because the correlation of Tschernitz et al. used the average of the positive and negative magnetic flux rather than the total unsigned flux as estimated by Schrijver et al. (2012), the reconnection flux for the 1947 region should be reduced by another factor of two from 3.0 × 1023 Mx to 1.5 × 1023 Mx (A. Veronig, personal communication, 2021), reducing the SXR class estimate to X180 (− 100, + 300) (1.4 (− 0.6, + 1.4) × 1033 erg) via Eq. (5). For the compromise total (average) unsigned flux of 4(2) × 1023 Mx for the April 1947 active region in Fig. 15 (with 50% flux involvement), Fig. 16b yields a SXR class of X80 (− 40, + 120) (with bolometric energy of 0.8 (− 0.3, + 0.7) × 1033 erg). For reasons given in 3.1.7(c) below, our preferred estimate based on Fig. 16b is the X180(− 100, + 300) value for an unsigned flux of 6 × 1023 Mx. Figure 17 gives an impression of the areas of sunspot groups that appear to be required to power flares of different magnitudes. The inset shows the large 1947 active region for comparison with a modelled spot group in the upper right of the figure that, based on the above calculation, could produce a 1033 erg flare.
Kazachenko et al. (2017) performed a similar study to that of Tschernitz et al. based on over 3000 flare events (C1.0-X5.4) by measuring the ribbons observed in the 1600 Å channel of SDO/AIA (Atmospheric Imaging Assembly; Lemen et al. 2012) instead of Hα spectroheliograms, and mapping these to SDO/HMI magnetograms. They find
$$ \Phi_{{\text{r}}} = 10^{22.13} { }\left( {\frac{{{\mathcal{C}}_{GOES} }}{{{\mathcal{C}}_{{GOES,{\text{ X}}1}} }}} \right)^{0.67} $$
(6)
Equation (6) was derived as a linear reduced major axis fit to logarithmic variables (Isobe et al. 1990). The slope of 0.58 derived by Tschernitz et al. (2018) for the fit line in Fig. 16b and that of Kazachenko et al. in Eq. (6) are statistically consistent. Toriumi et al. (2017) obtain a slope of 0.28 ± 0.10 for a sample of 51 events but for a small (M5.0-X5.4) range. For an active region with total unsigned flux of 4 × 1023 Mx and 50% flux involvement, Eq. (6) yields an ~ X55 SXR flare with radiative energy of ~ 6 × 1032 erg from Eq. (5), comparable to the ~ X45 and 5 × 1032 erg values inferred for the 1 September 1859 event.
(b) Estimates based on active region energy and released flare energy
As noted in Sect. 3.1.5, solar flares and eruptions are powered by some fraction of the excess or free (i.e., non-potential) magnetic energy (Efm) of an active region. Efm is distributed throughout the magnetic field of an active region. Sometimes, authors define a “free energy density” as if that would provide information on where the primary contributions to free energy would be located. However, despite the fact that the difference of the integral of the field energy for non-potential and potential fields can be mathematically written as the integral over the difference of what looks like a local quantity, namely
$$ E_{{\text{fm }}} = { }\mathop \int \limits_{V} \frac{{B_{{{\text{non}} - {\text{pot}}}}^{2} }}{8\pi }dV - { }\mathop \int \limits_{V} \frac{{B_{{{\text{pot}}}}^{2} }}{8\pi }dV = { }\mathop \int \limits_{V} \left( {\frac{{B_{{{\text{non}} - {\text{pot}}}}^{2} }}{8\pi } - \frac{{B_{{{\text{pot}}}}^{2} }}{8\pi }} \right)dV $$
(7)
this should not be interpreted to mean that the physical quantity of the resulting integrand maps to, say, the electrical currents that most contribute to the free energy. Efm is intrinsically a large-scale quantity.
Estimation of Efm is complicated by the facts that we cannot directly measure the coronal magnetic field and that models of that field based on the measurements of polarized photospheric light are, at best, uncertain at present. Nevertheless, as summarized by Emslie et al. (2012): “Numerous efforts have been undertaken to estimate nonpotential magnetic energies in active regions near disk center. The methods include: (1) using the magnetic virial theorem estimates from chromospheric vector magnetograms (Metcalf et al. 1995, 2005), (2) semi-empirical flux-rope modelling using Hα and EUV images with MDI LOS magnetograms (Bobra et al. 2008), and (3) MHD modelling (Metcalf et al. 1995; Jiao et al. 1997) and non-potential field extrapolation based upon photospheric vector magnetograms (Guo et al. 2008; Schrijver et al. 2008; Thalmann and Wiegelmann 2008; Thalmann et al. 2008). These methods are labor intensive, and uncertainties in their energy estimates are large. For example, error bars on virial free-energy estimates can exceed the potential magnetic energy. Also, there is considerable scatter in estimates from studies that employ several methods to analyze the same data (e.g., Schrijver et al. 2008). A couple of generalizations, however, can be made. Free energies determined by virial methods matched or exceeded the potential field energy, while free energies estimated using other techniques typically amounted to a few tens of percent of the potential field energy. Published values for free energies in analytic (Schrijver et al. 2006) and semi-empirical (Metcalf et al. 2008) fields meant to model solar fields also hover around a few tens of percent of the potential field energy.” Thus, as noted above, Emslie et al. (2012) adopted 30% of the modelled potential energy of an active region as their estimate of the free energy available for eruptive flares.
It may well be that a value of ~ 30% of the total magnetic energy of an active region represents a maximum value of the free energy beyond which active-region coronal fields cannot be stable. A model example of this can be found in the magnetohydrodynamic experiment performed by Aulanier et al. (2010, 2012) and analyzed in light of extreme flaring by Aulanier et al. (2013). In their bipolar region, one polarity is subjected to a rotational shear, which builds up until an instability in the field develops, manifested in the eruption of a flux rope mimicking a CME. Aulanier et al. (2013) note that the “CME itself was triggered by the ideal loss-of-equilibrium of a weakly twisted coronal flux rope …, corresponding to the torus instability (Kliem & Török 2006; Démoulin and Aulanier 2010)”. During the eruption, the field above their modelled bipolar region released 19% of the total magnetic energy, with most of that energy available for the thermal evolution of the corona (i.e., a flare) and the remaining 5% going into the bulk kinetic energy of the CME. While the fraction of total energy released in this numerical experiment approaches 30%, the ~ 20:1 apportionment between flare and CME is opposite to that of Emslie et al. (2012).
Aulanier et al. (2013) obtained relationships between active region length scale, typical field strength, flux, and energy content. Reformulation of their results yields the following scaling between the total energy E(AR) (erg) contained in the active region magnetic field, the region’s total unsigned flux \(\Phi \) (Mx), and the core field strength of the simulated spot pair, BC (Mx cm−2):
$$ E\left( {AR} \right) = 0.14 \left( {\frac{{B_{C} }}{{1000 {\text{G}}}}} \right)^{1/2} \Phi^{3/2} . $$
(8)
This equation permits an independent determination of the largest possible flare from that determined in Sect. 3.1.7(a). With the scaling of Eq. (8), a maximum field strength of BC = 3.5 kG (Aulanier et al. 2013) and an active region flux of Φmax = 4 × 1023 Mx yields a total active region energy E of 6.6 × 1034 erg. Apportioning per Sect. 3.1.5, the released energy (approximately one-tenth of active region magnetic potential energy) is 6.6 × 1033 erg, with bolometric energy (~ one-fourth of released energy) of ~ 1.7 × 1033 erg and a SXR class of ~ 240 via Eq. (5)Footnote 6.
In another estimate of this type, Toriumi et al. (2017) calculated the total energy released in a flare based on the integral of the flux contained in the flare ribbons over time. For an estimated total magnetic flux of 2 × 1023 Mx for the April 1947 active region (Fig. 15) and under the assumptions that the ratio of ribbon area to spot area was 0.85 and that two-thirds of the computed energy would be released in a flare, they obtained a released energy of ~ 1 × 1034 erg, which after Emslie et al. (2012) would translate to a bolometric energy ~ 2.5 × 1033 erg, and a ~ X410 SXR flare from Eq. (5).
The likely underestimation of CME energies noted in Sect. 3.1.5 will reduce the SXR class estimates based on the work of Aulanier et al. (2013) and Toriumi et al. (2017). Assuming a 6:1 (vs. 3:1) apportionment of released free energy between CMEs and flares, respectively, results in a SXR class estimate of X105 (X185) for Aulanier et al. (Toriumi et al.).
(c) Composite estimate
Altogether, the above estimates have a surprisingly low range of values, all lying within the X80 (− 40, + 120) range for the estimate from Fig. 16b, based on a total unsigned flux of 4 × 1023 Mx from Fig. 15. The lower limit SXR class and radiative energy estimates of Kazachenko et al. (X55; 6 × 1032 erg) and Tschernitz et al. (X40; 5 × 1032 erg) are comparable to the parameters for the 1 September 1859 and 4 November 2003 flares, which originated in spot groups less than half as large as April 1947. Because of this and the apparent upward curvature of magnetic flux with increasing group spot area in Fig. 15, we adopt: (1) the reconnection flux based estimate of X180 (− 100, + 300) given by Fig. 16b from Tschernitz et al. as our preferred estimate of the largest possible flare, and (2) the total unsigned flux value of 6 × 1023 Mx they used as the best estimate of the unsigned flux for the April 1947 active region. An unsigned flux of 6 × 1023 Mx would increase the SXR flare class (radiative energy) estimate of Kazachenko et al. to X105 (9.4 × 1032 erg), and those of Aulanier et al. and Toriumi et al. to ~ X240 (1.7 × 1033 erg) and ~ X4000 (1.3 × 1034 erg), respectively. We reject the ~ X4000 estimate of Toriumi et al. as a clear outlier—an overestimate apparently due to assumptions regarding spot:ribbon area ratio and the fraction of the computed energy released. From Fig. 8, a 10,000-year flare would have a bolometric energy of ~ 2 × 1033 erg versus the 1.3 × 1034 erg value calculated from Eq. (5) for the ~ X4000 estimate from Toriumi et al. Nominal SXR class and energy values of X180 and 1.4 × 1033 erg are approximately five and three times as large as the respective values (~ X40; ~ 5 × 1033 erg) for the 1859 and 2003 flares. Our preferred estimate of X180 (− 100, + 300), indicated by bold/italic font in Table 1, can be considered a composite because it encompasses the ~ X105 and ~ X240 estimates of Kazachenko et al. and Aulanier et al., respectively, with all three estimates based on an unsigned magnetic flux of 6 × 1023 Mx. The lower end of the X180 (− 100, + 300) estimate for the largest possible flare based on the April 1947 active region exceeds that of the largest flares observed since 1859 (from smaller spot groups) and the upper limit encompasses the X285 (± 140) SXR class estimated for the inferred flare for the 774–775 AD SEP event (Sect. 7.8).
Table 1 Estimates of the largest possible present-era solar flare based on an unsigned magnetic flux of 6 × 1023 Mx for the active region of April 1947 (peak daily area = 6132 µsh), and a 6:1 CME:flare energy apportionment Flares on Sun-like stars
Stellar flare research before Kepler
The first generally recognized observations of stellar flares (on the faint dwarf star L 726-8; also designated UV Ceti) occurred on 25 September 1948 (Joy and Humason 1949) and 7 December 1948 (Luyten 1949a, b). Luyten (1949a) noted that Edwin Carpenter’s photographic plates, taken on 7 December with the 36-in. telescope at Steward Observatory in Tucson (Fig. 18), give “the first observation of the extreme rapidity of the [flare brightness] change—to twelve times the original luminosity [of the star] in less than 3 min …” Precursor observations of stellar flares were made by Hertzsprung in 1924 (Bastian 1990; Gershberg 2005), Luyten (1926) and Van Maanen (1940, 1945). Subsequently, radio (Lovell 1963; Slee et al. 1963a, b; Orchiston 2004) and X- ray emissions (Heise et al. 1975; Güdel 2004) were found to accompany optical stellar flares.
Joy and Humason (1949) and Luyten (1949a, b) used the terms “flare-up” or “flare” when describing the stellar phenomenon but neither paper makes an analogy to solar flares. Luyten (1949b) suggested that the cause might be “the same as that which produces a nova, or an SS Cygni star.” In the preface of his monograph on solar activity in main sequence stars, Gershberg (2005) places the establishment of the similarity of solar and stellar flares in the mid-1960s, following the first high-time resolution spectra of flares on UV Ceti variables. As noted in Sect. 3.1.3, several lines of evidence support the close relationship of physical processes observed in stellar and solar flares.
Discovery of superflares on solar-type (vs. Sun-like) stars
X-ray observations show that young stars can have very large flares, termed “superflares” (Schaefer et al. 2000), defined as flares with radiative energy more than 1033 erg (~ 2–3 times more energetic than the largest directly observed solar flares; Emslie et al. 2012). These young stars rotate very fast with rotational periods of a few days, much less than solar sidereal rotation period of ~ 25 days at the equator. Generally, such fast rotating stars are strong X-ray sources even in their quiescent phase, indicating the presence of strong magnetic field and often with evidence that large fractions of the stellar surface are covered by starspots (Pevtsov et al. 2003). Consequently, it has been assumed that superflares would never occur on the present Sun, since the Sun is old and is slowly rotating.
However, by analyzing previously existing astronomical data, Schaefer et al. (2000) identified nine superflares with energies ~ 1033–38 erg in ordinary solar-type stars (i.e., G type main sequence stars with rotational velocities less than 10 km s−1; rotation period > 5 days). Concurrently, Rubenstein and Schaefer (2000) argued from analogy with RS CVn binary star systems that superflares on solar-type stars were caused by a “hot Jupiter” (Mayor and Queloz 1995; Marcy and Butler 1998; Schilling 1996), a Jovian-type exoplanet orbiting close-in to these stars. Rubenstein and Schaefer thus implied that—because of the lack of a hot Jupiter—the Sun is not a candidate to produce superflares. Cuntz et al. (2000) laid the framework for how a close-in giant planet could cause stellar activity via gravitational and magnetic interaction. Ip et al. (2004) used a numerical MHD simulation to show that a star-planet magnetic interconnection could lead to energy release comparable to that of a typical solar flare and Lanza (2008) explained the phase relation between stellar hot spots and planetary location in terms of such a linkage. Direct observational evidence between stars with close-in Jovian planets and superflares was lacking, however (e.g., Saar et al. 2004).
One of the authors of this review, Kazunari Shibata, questioned the hot-Jupiter hypothesis of Rubenstein and Schaefer (2000) because the RS CVn stars on which the hypothesis is based have short rotation periods due to tidal locking. Hence in 2010 Shibata started to encourage young researchers and students to search for superflares on solar-type stars (G type dwarfs) observed by Kepler. As a result, Maehara et al. (2012) reported 365 superflares (with radiative energy > 1033 erg) on 148 “solar-type” stars (defined as G-type main sequence stars with effective temperature of 5100–6000 K and log (surface gravity (cm/s2) ≥ 4.0) during 120 days of Kepler observations. For more restrictive criteria, they identified [9] superflares on 5 “Sun-like” stars in the sample with effective temperatures of 5600–6000 K and rotational periods between 11.0 d and 17.1 d (Maehara et al.; Supplementary Information). As a reflection of Shibata’s encouragement, the Nature paper reporting these results included five undergraduate researchers as co-authors (T. Shibayama, S. Notsu, Y. Notsu, T. Nagao, and S. Kusaba).
Later, Shibayama et al. (2013) extended the survey and confirmed the work of Maehara et al. by finding 1547 superflares (including those reported by Maehara et al. 2012) on 279 G-type dwarf stars during 500 days (from 2009 May to 2010 September) of Kepler observations. In all, they found 44 superflares on 19 Sun-like stars. Of these 19 Sun-like stars, three had rotational periods longer than the solar rotation period of ~ 25 days, suggesting that slowly rotating stars like the Sun could exhibit superflares.
Maehara et al. (2012) and Shibayama et al. (2013) argued that the superflares they discovered were not due to hot Jupiters orbiting the superflare stars. Quoting from Shibayama et al., “According to the Kepler candidate planet data explorer (Batalha et al. 2013), 2321 planets have been found in 1790 stars among 156,453 stars. Hence, the probability of finding exoplanets orbiting stars is about 1%. Howard et al. (2012) showed that the probability of finding a hot Jupiter was 0.5%. However, none of our superflare stars (279 G-type dwarfs [of which 69 had rotation periods > 10 d]) have a hot Jupiter according to the data explorer. For a G-type dwarf with a hot Jupiter, the probability of [detecting] a transit of the planet across the star is about 10% averaged over all possible orbital inclinations (Kane and von Braun 2008). If all of our 279 [69] superflare stars are caused by a hot Jupiter as suggested by Rubenstein and Schaefer (2000), Kepler should detect 28 [7] of them from transits” versus the zero that were found. Recently, for a sample of 265 solar-type super flare stars (G-type main-sequence; effective temperature of 5100–6000 K; 139 with rotation periods > 10 days), with 43 in common with the 279 G-type dwarfs of Shibayama et al. 2013), Okamoto et al. (2021) found only three stars with candidate exoplanets in the NASA Exoplanet Archive (https://exoplanetarchive.ipac.caltech.edu/).
Figure 19a shows a typical example of a superflare observed by Kepler, with a spike-like increase peaking at 1.5% of the stellar brightness for a few hours. Flares on Sun-like (solar-type) stars have brightness variations in the range of ~ 0.1 to ~ 20% of the stellar luminosity (Notsu et al. 2013a). In contrast, one of the largest solar flares in the past 20 years (X17 SXR flare on 28 October 2003 shown in Fig. 13; Woods et al. 2004; Kopp 2016) showed only a ~ 0.03% solar brightness increase for ~ 20 min (FWHM). The estimated bolometric energy of the superflare in Fig. 19a is ~ 1035 erg, ~ 200 times larger than that for the largest solar flare at ~ 5 × 1032 erg.
In Fig. 19a the stellar brightness itself shows significant rotational variation with an amplitude of ~ 1% at a characteristic time of 10–15 days. Almost all superflare stars show such a variation ranging from 0.1 to 1% (Maehara et al. 2012) versus ~ 0.1% for the 11-year solar cycle TSI variation (Willson and Hudson 1991; Kopp 2016). This stellar brightness variation can be interpreted in terms of the rotation of a star with a substantial coverage by unevenly distributed starspots (Strassmeier et al. 2002, and references therein; Strassmeier 2009; Notsu et al. 2013a; Namekata et al. 2019), such as depicted in Fig. 19b. The interpretation of Kepler brightness variations in terms of rotation of large area starspot groups can be used to measure the rotation period of stars and to infer the fractional coverage of starspots (technically, one can only infer the maximum difference between opposing hemispheres of total spot coverage; see below), or total magnetic flux, assuming the average magnetic flux density is the same as that of a large sunspot. The derivation of superflare star, and superflare, parameters from Kepler light curves will be discussed in Sect. 3.2.4 below.
Before proceeding further, we need to add a caveat. Up to this point, we have been using the terms radiative energy, bolometric energy, and TSI interchangeably, taking it to mean flare-radiated energy time-integrated over all wavelengths. For Kepler white light flare observations, the term “bolometric flare energy” has a different, more limited, definition. The “bolometric energy” as referenced for Kepler here is the flare-radiated energy in the blackbody continuum, assuming the flare to have an effective temperature of 10,000 K to convert the emission that occurs in the Kepler 4000–9000 Å bandpass into a bolometic energy (e.g., Notsu et al. 2013a) based on a solar study by Kretzschmar (2011). We further note that there is no consideration of line versus continuum emission in the photosphere and chromosphere or allowance for coronal emission. Thus far such computations have not been done for solar-like stars other than the Sun. Consequently, the Kepler flare bolometric energies could be underestimated by up to a factor of ten in comparison with solar values based on actual TSI measurements (Osten and Wolk 2015).
Verifying the Sun-like nature of superflare stars
(a) Initial verification
It is long known that close binary systems such as the RS CVn-type (Hall 1976) can produce superflares with energies ~ 1035 erg or more (e.g., Doyle et al. 1991). Doppler observations can be used to eliminate such binary systems (except for systems seen (nearly) pole-on), as well as single stars rotating much faster than the Sun, from the Kepler sample of nominally Sun-like superflare stars. Other spectroscopic observations can be used to substantiate the Sun-like character of superflare stars. These include measurements of: Hα and Ca II 8542 fluxes (gauges of chromospheric activity; Linsky et al. 1979; Herbig 1985; Foing et al. 1989; Soderblom et al. 1993), Ca II K flux (indicator of the stellar mean magnetic field strength; Skumanich et al. 1975; Schrijver et al. 1989); Fe/H ratio (metallicity) and Li abundance (both indicators of stellar age; Edvardsson et al. 1993; Skumanich 1972; Soderblom 2010), and equivalent widths of Fe I and Fe II lines (effective temperature and surface gravity; Takeda et al. 2002, 2005; Valenti and Fischer 2005). Early spectroscopic studies of superflare stars were conducted by Notsu et al. (2013b; one star), Nogami et al. (2014; two stars), and Wichmann et al. (2014; 11 stars). The detailed discussion of the 11 stars examined in Wichmann et al. illustrates the complexity of the task, with all indicators of “Sun-likeness” seldom pointing in the same direction.
Notsu et al. (2015a, b) reported on their spectroscopic investigation of 34 non-binary solar-type (G-type main sequence) superflare stars observed with the High Dispersion Spectrograph (HDS; Noguchi et al. 2002) on the Japanese 8.2-m Subaru telescope in Hawaii. As a check on light-curve based stellar rotation periods (and inferred rotation rates), Notsu et al. (2015a), used HDS data to determine the projected rotation rates (v sin i; where i is the inclination angle between the stellar rotational axis and the observer’s line of sight, measured from the pole; e.g., for i = 0°, the stellar rotation axis points directly along the line-of-sight) of these 34 starsFootnote 7 and parameters such as metallicity, effective temperature, and surface gravity (all via Fe I and Fe II line widths) to isolate a sample of seven Sun-like stars that met the following criteria: 5600 ≤ Teff ≤ 6000 K, log g ≥ 4.0, and rotation periods (based on brightness variation) > 10 days.
The seven Sun-like stars identified by Notsu et al. (2015a) are listed in the first seven rows of Table 2 along with their effective temperature (Teff), surface gravity (g), amplitude of brightness variation (BV), metallicity (Fe/H), mean magnetic flux density (fB), spot area (Aspot), rotational velocity (determined both photometrically (vlc (lc = light curve)) and spectroscopically (v sin(i)), rotation period (Prot), and bolometric energy of the largest observed flare (Max flare). Recently an additional Sun-like star (KIC 11253827) has been identified by Notsu et al. (2019). Its parameters are given in row 8. Corresponding parameters for the Sun are given in the bottom row.
Table 2 Sun-like superflare stars identified by Kepler (from Nogami et al. 2014; Notsu et al. 2015a, b, 2019) Light curves for the two stars in Table 2 that were analyzed in detail by Nogami et al. (2014) are shown in Fig. 20. Those authors reported photometrically-based rotation periods of 21.8 days and 25.3 days for KIC 9766237 and KIC 9944137, respectively. Subsequently, these two periods were reduced in Notsu et al. (2015a, see their Appendix A1 (Supplementary data)) to the 14.2 day and 12.6 day values shown in Table 2. The Nogami et al. (2014) periods were calculated by Shibayama et al. (2013) from Kepler Quarters 0–6 data while those from Notsu et al. (2015b) were based on updated Quarter 2–16 data.
Witzke et al. (2020) document the difficulty of obtaining photometry-based light curves for “solar-like” stars by analyzing a set of synthesized light curves for stars that are identical to the Sun except for their metallicity. They find a minimum in the detection rate of the rotation period for stars having a metallicity value close to that (0.01) of the Sun because of the close balance of facular and spot contributions in our star. Witzke et al. substantiate this result by considering a sample of Kepler solar-like stars with a range of metallicities from − 0.35 to 0.35 and near-solar effective temperatures and photometric variabilities. The underlying assumption of long-term coherence, i.e., of long-lived spots comparable to the stellar rotation period, is not tenable for the Sun (in its present state). As of 2019, KIC 9766237 and KIC 9944137 appeared to be the most Sun-like superflare stars in terms of their brightness variation (BV) and rotation velocities (Y. Notsu, personal communication, 2019), but as noted below, even the reduced rotation periods from Notsu et al. (2015a) are now in question. The two associated superflares (one per star) had energies of ~ 1034 erg, at the lower end of the Kepler superflare range and had associated Sun-like spot areas of ≲ 1000. Such spot areas are at or below the limit of Kepler detection (Nogami et al. 2014), but see Sect. 3.2.4 below for a different interpretation of these inferred values.
Maehara et al. (2017) identified two other candidates for Table 2, both with photometrically-determined rotation periods > 20 days. KIC 6347656 and KIC 10011070 have spot area brightness-variation-based rotation periods (spot areas, Teff) of 28.4 days (3600 µsh, 5623 K) and 24.0 days (3400 µsh, 5669 K), respectively. Each of these stars had one superflare with bolometric energy of 3.1 × 1034 erg for KIC 6347656 and 2.6 × 1034 erg for KIC 10011070. The two stars have yet to be examined spectroscopically (Y. Notsu, personal communication, 2019).
In Table 2, the mean magnetic flux densities (fB) for the Sun-like stars cannot be distinguished from that of the Sun. The eight listed stars in Table 2 have a median flux density of 4(± 24) G versus ~ 10 G for the Sun. The values of < fB > are derived from a calibration against the intensity of the stellar Ca II infrared triplet relative to its local continuum (Notsu et al. 2015b). However, this calibration assumes observations of an area around a single active region. The authors do not discuss the consequences of using such observations to calibrate the signal against the disk-integrated signal. Furthermore, the uncertainties that they list are characteristic of the spread of points around the mean relationship, not the uncertainty in the mean itself. Finally, we note that that relationship when applied to the Sun yields a mean magnetic flux density of 0.2G, where observations show values of ~ 10 G, i.e., ~ 1.5 orders of magnitude larger. In addition, Notsu et al. (2015a, b) make no mention of center-to-limb effects for the transformation from their single spot group measurement to a disk-integrated signals for stars (cf. Namekata et al. 2017).
(b) A subsequent challenge
Subsequent development and refinement of the Kepler data base challenged the case for superflares on Sun-like stars. Quoting from Notsu et al. (2019): “The recent Gaia-DR2 stellar radius data (Berger et al. 2018; Lindegren et al. 2007, 2018) have suggested the possibility of severe contaminations of subgiant stars in the classification of Kepler solar type (G-type main-sequence) stars used for the statistical studies of superflares … The classification of solar-type stars in our previous studies … was based on Teff and log g values from the Kepler Input Catalog … and there can be large differences between the real and catalog values. For example, Brown et al. (2011) reported that uncertainties of Teff and log g in the initial KIC are ± 200 K and 0.4 dex [decimal exponent], respectively. … In a strict sense, we cannot even deny an extreme possibility that all of the slowly rotating Sun-like superflare stars [we considered] were the results of contaminations of subgiant stars. In addition, the Kepler solar-type superflare stars discussed in our previous studies can include some number of binary stars [excepting the events in Table 2 that were confirmed to be single, and Sun-like, by spectroscopic examination]. This is a problem, especially for investigating whether even truly Sun-like stars can have large super flares or not.“
To address the above concern, Notsu et al. (2019) revisited earlier analyses (e.g., Maehara et al. 2012, 2015, 2017; Shibayama et al. 2013) by using the Gaia-DR2 stellar radius estimates (Berger et al. 2018) and Teff values updated in DR25-KSPC (Mathur et al. 2017; Pinsonneault et al. 2012). In so doing, they found that of the 245 (of 279) stars in the Shibayama et al. (2013) sample of solar-type (G-type dwarfs) superflare stars (from 2009 May to 2010 September) for which Berger et al. (2018) determined a radius, 108 (~ 45%) were subgiants (with one red giant) and, of the remaining 136 identified as main-sequence stars, only 106 had Teff,DR25 from 5100–6000 K. Spectroscopic analysis (3200–10,000 Å) was then conducted using the Apache Point Observatory 3.5 m telescope for 18 main-sequence superflare stars observed by Kepler at 1-min cadence for which Gaia-DR2 (Release 2) stellar radius and Teff,DR25 values were available; 13 of these 18 stars were found to be single.
(c) Further development
In a further development based on the entire Kepler primary mission data set (Kepler Data Release 25 (Thompson et al. 2016; Mathur et al. 2017); 2009 May to 2013 May) and the Gaia-DR2 catalog, Okamoto et al. (2021) used a flare detection method that employed high-pass filtering to remove low-frequency rotational variations due to starspots to identify superflares on 15 Sun-like stars (based on Teff from 5600 to 6000 K, a more stringent Prot requirement of > 20 d (with Prot taken from McQuillan, et al. 2014), and age of 4.6 Gyr). The parameters for these 15 Sun-like superflare stars are given in Table 3 (for 15 vs. the 16 in the table, see Flag 2). This list includes the two Sun-like stars with rotation periods > 20 days identified by Maehara et al. (2017) noted above. The range of spot areas from ~ 1500 to 13,000 µsh for the superflare stars in Table 3 are comparable to those in Table 2 while, as noted by Okamoto et al. (2021), the faster rotating stars in Table 2 are capable of producing stronger superflares (by up to an order of magnitude) for the listed events. At present the case for superflares on Sun-like stars is based largely on Table 3.
Table 3 Parameters of Sun-like superflare stars In a reappraisal of KIC 9766237 and KIC 9944137 (from Table 2) for which light curves are shown in Fig. 20 above, Okamoto et al. (2021) assess them to be “Sun-like star candidates”, rather than Sun-like stars, because “the amplitude of the rotational brightness variation is so small that [its] value is not reported in McQuillan et al. (2014), and the accuracy of the Prot value is considered to be low.” In contrast, as discussed in Sect. 8.3.1 below, the inability to photometrically detect a rotation period may be a signature of a Sun-like star.
Inferring superflare star properties from brightness variations
From the Kepler photometric observations, it is possible to estimate the stellar rotation period, starspot area, and flare energy. We will consider each of these in turn.
Maehara et al. (2012) and Shibayama et al. (2013) calculated the rotational periods of superflare stars by discrete Fourier analysis of the stellar brightness variations for the first 500 days of the Kepler mission. They chose the largest peak that exceeded the error level as the stellar rotation period. Alternatively, McQuillan et al. (2014) used an automated autocorrelation function (ACF) approach to determine rotation periods based on the entire Kepler data set.
The lower limit of the spot area of superflare stars (Aspot) can be inferred from the ΔFrot (defined as the full amplitude of the rotational brightness variation normalized by the average brightness) as follows:
$$ \Delta F_{{{\text{rot}}}} \approx { }\left[ {1 - \left( {\frac{{T_{{{\text{spot}}}} }}{{T_{{{\text{star}}}} }}} \right)^{4} } \right]\frac{{A_{{{\text{spot}}}} }}{{A_{{{\text{star}}}} }} $$
(9)
where Tspot is the spot temperature, Tstar is the effective temperature of the unspotted photosphere of the star, and Astar is the apparent area of the star (Maehara et al. 2012, 2017). The spot area in this equation is a lower limit because it is assumed that there is a reference hemisphere that has no (or negligible) starspot coverage (e.g., Fig. 3b in Notsu et al. 2013a). Notsu et al. (2019) used a relationship between star temperature and spot temperature (Maehara et al. 2017) for which Tstar (in the range from 5600 to 6000 K) was taken from the Kepler catalog and Tspot was determined spectroscopically from intensities of the photospheric lines of sunspots.
Figure 21 (taken from Maehara et al. 2017) is a downward cumulative frequency distribution (solid line) showing the occurrence rate of active regions of a given spot area for solar-type stars (defined to be early G- and late F-type main-sequence stars that meet both of the following criteria: (1) 5600 K < Teff < 6300 K and log g > 4.0 in Huber et al. (2014); and (2) 5100 K < Teff < 6000 K and log g > 4.0 in Brown et al. (2011)). The corresponding solar distribution (dotted line) is shown for ~ 140 years of solar data. The continuation of the extrapolated power-law fit to the distribution for solar-type stars (dashed line) through that of solar active regions suggests that both sunspots and larger starspots might be produced by the same physical process and that the Sun is capable of producing spot groups much larger than observed in modern times. Maehara et al. (2017) suggest that the steep drop-off of the solar distribution is due to “the lack of a ‘super-active’ phase on our Sun during the last 140 years”. This seems problematic, however, because the four solar cycle interval of solar activity from ~ 1945 to 1995, termed the modern grand maximum, is considered to be one of the strongest in the past 2000–4000 years (Solanki et al. 2004; Usoskin et al. 2006a; Muscheler et al. 2006; Usoskin 2017). Figure 21 indicates that a ~ 30,000 µsh spot region, region ~ 5 times larger than that of April 1947, would be expected to be observed once every ~ 1000 years, with no evidence of a region approaching this size over the ~ 400 years of telescopic observation of the Sun. That said, the time interval for which we have detailed knowledge of the Sun's variability is minuscule compared to the time-scale of stellar evolution. As discussed in Sect. 8.3.1 below, recent analysis of Kepler photometric and Gaia astrometric data by Reinhold et al. (2020) raises the possibility that our Sun is currently in a state of subdued activity in comparison with the bulk of Sun-like stars.
After Shibata et al. (2013), the energy (in erg) of a stellar flare (Eflare) in a sunspot group spot with magnetic flux density B, scale-length L, and area Aspot has an upper limit determined by the total magnetic energy stored in a volume A3/2 near the spot, i.e.,
$$ E_{{\text{flare }}} \approx fE_{{{\text{mag}}}} \approx f\frac{{B^{2} L^{3} }}{8\pi } \approx f\frac{{B^{2} }}{8\pi }{ }A_{{{\text{spot}}}}^{3/2} \approx 7 \times 10^{32} { }\left( {\frac{f}{0.1}} \right)\left( {\frac{B}{{10^{3} G}}} \right)^{2} \left( {\frac{{A_{{{\text{spot}}}} }}{{3 \times 10^{19} {\text{ cm}}^{2} }}} \right)^{3/2} { } \approx 7 \times 10^{32} {\text {(erg)}} { }\left( {\frac{f}{0.1}} \right)\left( {\frac{B}{{10^{3} G}}} \right)^{2} \left( {\frac{{A_{{{\text{spot}}}} /\left( {2\pi R^{2}_{ \odot } } \right) }}{0.001}} \right)^{3/2} $$
(10)
where f (0.1; Sect. 3.1.5) is the fraction of magnetic energy that is released during the flare and R is the solar radius (to which the stellar radius is scaled). To scale the flare energy determined by Eq. (10) to SXR class, Shibata et al. (2013) assumed linear proportionality between these parameters from consideration of previous observational estimates of energies of typical solar flares (e.g., Benz 2008; also used by Gopalswamy et al. 2018). The linear scaling of Shibata et al. (2013) differs from that of Schrijver et al. (2012; Eq. (5) in this paper with an exponent of 0.72) as well as the similar scaling in Tschernitz et al. (2018; 0.79), both of which are based on the data in Kretzschmar (2011). Namekata et al. (2017) show that an exponent of 1.0 lies between those for an ordinary least squares (OLS) fit (0.84 ± 0.04) and a linear regression bisector fit (1.18 ± 0.04; Isobe et al. 1990) to a scatter plot of the logs of HMI white-light flare energy versus SXR peak flux for a sample of M- and X-class flares. It is not clear that the bisector fit is more appropriate in this case than the OLS fits in log–log space used by Schrijver et al., Tschernitz et al., and Namekata et al., which yielded slopes from 0.72 to 0.84. The Schrijver et al., Tschernitz et al., and Shibata et al. scalings are shown in Fig. 22 along with the Kretzschmar data points. The Tschernitz et al. and Schrijver et al. fits differ slightly because Schrijver did not correlate bolometric energy and GOES SXR class directly but first related TSI to SXR fluence and then converted SXR fluence to SXR class after Veronig et al. (2002a). The three scalings in Fig. 22 agree reasonably well for events in the X10-X1000 range of interest here, although Shibayama et al. (2013; their Appendix B) note an apparent disagreement between the observed frequency of nanoflares and the expected rate based on Kretzschmar's relation (Eq. (3) above) between bolometric energy and SXR energy.
Figure 23 from Okamoto et al. (2021) is a combined scatter plot of (1) the SXR classes of solar flares (right-hand y-axis) and (2) the bolometric energies of Kepler superflares on the 15 solar-type stars in Table 3 (left-hand axis) versus their spot areas. For both sets of flares, the energy to flare class scaling of Shibata et al. (2013) in Fig. 22 is used to infer the missing parameter—bolometric energy for solar flares and SXR class for stellar superflares. The SXR flare versus spot area comparison was patterned after Sammis et al. (2000).Footnote 8 For the stellar flare data points, spot areas were inferred from Eq. (9) with bolometric energy given by Kepler photometry. The solid and dashed slanting lines given by Eq. (10) represent the released flare energy (with f = 0.1) for a given de-projected spot area (inclination angle i assumed to be 90°) on the x-axis and spot field strengths of 3 kG and 1 kG, respectively.
The apparent correlation in the peak SXR flare flux versus solar spot group area in Fig. 23 has been discounted by Hudson (2021) because of “a strong selection bias, resulting from the under-reporting of weaker GOES events due to higher background levels during active times (Wheatland 2010)”. This selection effect will also apply to the flares observed by Kepler, for which the detection completeness of 1033 erg flares is 0.001 (Maehara et al. 2012). That said, the increase in the largest observed flare energies of stellar flares relative to those of solar flares in Fig. 23, in correspondence with the larger maximum spot areas inferred for stellar flares, is in accord with expectations—particularly because the starspot areas are likely underestimated due to the “null hemisphere” hypothesis. From this figure, Okamoto et al. (2021) conclude that “stellar superflares [represent] the magnetic energy release stored around the starspots, and the process is the same as that of solar flares”.
Occurrence frequency of superflares on Sun-like stars
Figure 24 (adapted from Okamoto et al. 2021; see also Schrijver et al. 2012; Maehara et al. 2015; Lin et al. 2019; Notsu et al. 2019) shows the unified figure of occurrence frequency of flares as a function of flare energy, for solar flares and superflares on Sun-like stars (based on short time cadence Kepler data). There is an apparent smooth distribution from the smallest EUV solar flares all the way to the largest bolometric flares on Sun-like stars. It is remarkable that superflare frequency is roughly on the same line as that for solar flares, microflares, and nanoflares, dN/dE ~ E−1.8, suggesting the same physical mechanism for both solar and stellar flares (see Sect. 3.1.3). In an early result for a sample of solar-type stars observed by the Transiting Exoplanet Survey Satellite (TESS; Ricker 2015), Tu et al. (2020) reported a slightly steeper (− 2.16 ± 0.10) slope, but within the uncertainties of the Kepler results. Recently, Aschwanden and Güdel (2021) obtained a slope of − 1.82 ± 0.005 for the fluence distribution of 162,262 flares on stars of types A-M and giants observed by Kepler at optical wavelengths (Yang and Liu 2019) veraus a slope of − 1.98 ± 0.11 for the (background- subtracted) peak fluxes of 338,661 GOES SXR flares (Aschwanden and Freeland 2012).
A power-law slope of ~ − 1.8 pertains to a wide variety of solar emissions (Aschwanden et al. 2016) and is in theoretical agreement, within uncertainties, with the − 1.67 value for hard X-ray (> 10 keV) peak fluxes obtained by Aschwanden (2012) from the SOC model for solar flares (Lu and Hamilton 1991). The similarity of slopes in these various studies may reflect the universality of the reconnection process in converting magnetic energy to radiative energy via particle acceleration in different types of settings (the Sun and other stars; Benz and Güdel 2010; Aschwanden and Güdel 2021) as well as in different types of flares (e.g., confined and eruptive; Harra et al. 2016; Tschernitz et al. 2018; cf., Cliver and D'Huys 2018). The Neupert effect implies a power-law slope of − 2.0 for solar SXR flares (Aschwanden and Freeland 2012).
In their earlier analysis of a power-law distribution of Sun-like superflare stars (defined as G-type main sequence stars with 5800 ≤ Teff < 6300 and rotation period > 10 days) observed during the first 500 days of the Kepler mission, Maehara et al. (2015) determined that the occurrence frequency of 1033 erg solar flares would be once every 500 years. Subsequently, Notsu et al. deduced that flares with energy of > 1034 to ≲ 5 × 1034 erg occur on old (~ 4.6 Gyr), slowly rotating (Prot ∼ 25 days) Sun-like stars approximately once every 2000–3000 years. For a larger sample size than that of Notsu et al. and using a gyrochronology correction, Okamoto et al. (2021) inferred that solar superflares with energies of 1034 erg (SXR class = X1000) could occur on slowly rotating Sun-like stars once every ~ 6000 years. This rate is preliminary, pending spectroscopic analysis of metallicity, surface gravity, binarity, and rotation period of the 15 (presumably) Sun-like stars in Table 3 on which it is based (see Appendix B in Okamoto et al.). In addition, the possible underestimation of the energy of Kepler superflares (Osten and Wolk 2015) noted in Sect. 3.2.2 will need to be addressed.
Five of the 15 Sun-like stars in Table 3 produced more than one such superflare during the ~ 4 years of Kepler observations that were considered. Statistically, this seems incongruent with ~ 1034 erg superflares happening only once in six millennia. This might suggest clustering of superflares on Sun-like stars (as witnessed for intense solar flares; e.g., Cliver et al. 2020b) but it also might mean that the sample of candidate Sun-like stars in Table 3 is contaminated by faster rotators (a possibility noted for five stars with flag = 1 in Table 3; see Appendix A in Okamoto et al.), underscoring the need for spectroscopic analysis. Until that is done, it is uncertain whether the stars in Table 3 are truly Sun-like. As Okamoto et al. (2021) explicitly note, “the number of Sun-like superflare stars (Teff = 5600–6000 K, Prot ∼ 25 days, and t ∼ 4.6 Gyr) that have been investigated spectroscopically and confirmed to be “single” Sun-like stars, is now zero …” Thus the question posed in the title of Notsu et al. (2019)—Do Kepler Superflare Stars Really Include Slowly Rotating [with rotation periods ~ 25 days] Sun-like Stars?—remains unanswered. Until it is answered affirmatively, the Kepler observations of superflares (≥ 1033 erg) can only hint at the possibility that the Sun is capable of producing such flares.
While further spectroscopic work is required to substantiate the implications of superflares on Sun-like stars for extreme solar activity, observations of solar flares with bolometric energies within a half-decade of 1033 erg (Sect. 3.1.2 above) suggest that the Sun is capable of producing a threshold level superflare over time, while analyses of historical solar proton events based on cosmogenic radionuclides (Sect. 7 below) provide evidence that it has already done so.