Observations of white-light flares in NOAA active region 11515: high occurrence rate and relationship with magnetic transients
aa r X i v : . [ a s t r o - ph . S R ] J a n Astronomy&Astrophysicsmanuscript no. ms c (cid:13)
ESO 2018June 13, 2018
Observations of white-light flares in NOAA active region 11515:high occurrence rate and relationship with magnetic transients
Y. L. Song , , H. Tian , M. Zhang , , and M. D. Ding School of Earth and Space Sciences, Peking University, Beijing 100871, China;e-mail: [email protected] State Key Laboratory of Space Weather, Chinese Academy of Sciences, Beijing 100190, China Key Laboratory of Solar Activity, National Astronomical Observatories, Chinese Academy of Sciences, Beijing 100012, China School of Astronomy and Space Science, University of Chinese Academy of Sciences, Beijing 100049, China School of Astronomy and Space Science, Nanjing University, Nanjing 210093, China
ABSTRACT
Aims.
There are two goals in this study. One is to investigate how frequently white-light flares (WLFs) occur in a flare-productiveactive region (NOAA active region 11515). The other is to investigate the relationship between WLFs and magnetic transients (MTs).
Methods.
We use the high-cadence (45s) full-disk continuum filtergrams and line-of-sight magnetograms taken by the Helioseismicand Magnetic Imager (HMI) on board the Solar Dynamics Observatory (SDO) to identify WLFs and MTs, respectively. Imagestaken by the Atmospheric Imaging Assembly (AIA) on board SDO are also used to show the morphology of the flares in the upperatmosphere.
Results.
We found at least 20 WLFs out of a total of 70 flares above C class (28.6%) in NOAA active region 11515 during its passageacross the solar disk ( E ◦ ∼ W ◦ ). Each of these WLFs occurred in a small region, with a short duration of about 5 minutes. Theenhancement of white-light continuum intensity is usually small, with an average enhancement of 8.1%. The 20 WLFs observed werefound along an unusual configuration of the magnetic field characterized by a narrow ribbon of negative field. Furthermore, the WLFswere found to be accompanied by MTs, with radical changes in magnetic field strength (or even a sign reversal) observed during theflare. In contrast, there is no obvious signature of MTs in those 50 flares without white-light enhancements. Conclusions.
Our results suggest that WLFs occur much more frequently than what was previously thought, with most WLFs beingfairly weak enhancements. This may explain why WLFs are not frequently reported. Our observations also suggest that MTs andWLFs are closely related and appear co-spatial and co-temporal, when considering HMI data. A larger enhancement of WL emission isoften accompanied by a larger change of the line-of-sight component of the unsigned magnetic field. Considering the close relationshipbetween MTs and WLFs, many previously reported flares with MTs may be WLFs.
Key words.
Sun: activity-Sun: chromosphere-Sun: photosphere-Sun: flares-Sun: magnetic fields
1. Introduction
White-light flares (WLFs) are defined as flares with a sudden en-hancement of emission in the optical continuum ( ˇSvestka 1970;Neidig 1989), which are often believed to be very rare ( ˇSvestka1966; Neidig & Cliver 1983; Fang et al. 2013). Since the ob-servation of the first WLF in 1859 (Carrington 1859, Hodgson1859), the number of WLFs recorded in literature is very smallcompared to the total number of solar flares. Only about 150WLFs have been reported conclusively in the literature, up to thebeginning of this century (Fang et al. 2013). Though the num-ber of observed WLFs is small, they are important because theychallenge our knowledge on the transportation of flare energy(Neidig 1989) and the heating mechanisms of the lower solaratmosphere (Ding et al. 1999a). To understand WLFs, severalheating mechanisms have been proposed: electron beam bom-bardment (Hudson 1972, Aboudarham & H´enoux 1986), soft-X-ray irradiation (H´enoux & Nakagawa 1977), Alfv´en wave dis-sipation (Emslie & Sturrock 1982, Fletcher & Hudson 2008),backwarming (Machado et al. 1989, Metcalf et al. 1990, Heinzel& Kleint 2014) and chromospheric condensation (Gan et al.1994, Kowalski et al. 2015b).
Send o ff print requests to : H. Tian Based on di ff erent characteristics in observations, WLFs areclassified into two types (Machado et al. 1986). For type I WLFs,there is a strong correlation between the time of white-light en-hancement and the peak time of hard X-ray (HXR) and mi-crowave radiations. Also the Balmer lines are often very broadand strong (Fang & Ding 1995). For type II WLFs, there areno such characteristics (Ding et al. 1999a, 1999b). This classi-fication implies that WLFs may have di ff erent origins of WLemission and di ff erent heating mechanisms.Temporal and spatial relationship between the enhancementsof white light, HXR and radio emissions has been frequently in-vestigated. WL kernels are usually found to be co-spatial withhard X-ray sources (Hudson et al. 1992; Metcalf et al. 2003;Chen & Ding 2005, 2006; Krucker et al. 2011; Hao et al. 2012;Cheng et al. 2015; Kuhar et al. 2016; Yurchyshyn et al. 2017).Mart´ınez Oliveros et al. (2012) calculated the centroidal heightsof the HXR source and WL emission for a flare close to the so-lar limb, and found that the mean heights above the photosphereare 305 ±
170 km and 195 ±
70 km, respectively. Krucker et al.(2015) studied three WLFs at the solar limb and found that thecentroids of WL and HXR ( ≥
30 keV) sources share a similarheight, which is about 300 −
450 km above the limb. Watanabeet al. (2010) studied an X1.5 WLF and found that the electronacceleration is closely correlated to the white-light production in time, space and power. Recently, Huang et al. (2016) studied25 stronger flares including 13 WLFs and found that the popula-tion of high energy electrons is larger when the WL emissionis stronger. Kleint et al. (2016) investigated an X1 WLF andfound that the energy deposited by electrons was su ffi cient forthe extra ultraviolet (UV) and visible continuum emission. Leeet al (2017) analyzed an X1.6 WLF and suggested that the WLemmision enhancement was directly produced by non-thermalelectrons. Based on a statistical analysis of 43 WLFs, Kuhar etal. (2016) found that the electrons of 50 keV are the main en-ergy source for WL emission. All these reported WLFs are X-or M-class flares.Comparatively, enhanced WL emission has been less fre-quently observed in smaller flares. However, with the increasingsensitivity of detectors, more and more small WLFs have beendiscovered. Hudson et. al (2006) studied 11 WLFs including 4C-class flares using observations obtained with the TransitionRegion and Coronal Explorer ( TRACE ; Handy et al. 1999) andfound that the minimum enhancement of WL emission is about8%. Another example is given by Jess et al.(2008) who observeda C2.0 WLF. The duration of the WL emission is about 2 min-utes and the diameter of the white-light kernel is less than 0 . ′′ .The enhancement of WL emission is above 300%. In these cases,the WL sources occupy only a very small fraction of area of theentire flare ribbons. At present, there are many questions thatare unclear and require further study, such as, how these smallWLFs are produced, whether they belong to type-I or type-IIWLFs, and what special conditions in the WL sources are re-quired compared to the non-WL flare ribbons.Although many more WLFs have been discovered in re-cent years than before, it is still unknown how frequently WLFsoccur. In previous studies, WLFs are casually captured by vi-sual inspection from either images or spectra. Doing so in-evitably misses some WLFs with weak intensities and smallsizes. Therefore, to explore the above question, a systematic sur-vey of WL emission for an active region with continuous obser-vations of a few days is required.Magnetic field changes associated with solar flares have beenreported in many previous studies (e.g., Severny 1964; Tanaka1978; Patterson 1984; Chen et al. 1989; Kosovichev & Zharkova2001). These changes are usually classified into two categories.One is flare-associated rapid and permanent changes which arethought to be real changes of magnetic field, due to the fact thatthe magnetic fields become more horizontal at the polarity in-version lines of flaring regions (Hudson et al. 2008; Wang &Liu 2010; Fisher et al. 2012). The other is the so-called mag-netic transients (MTs; e.g., Kosovichev & Zharkova 2001; Qiu& Gary 2003; Zhao et al. 2009), which are usually believed to bean observational artifact produced by the changes of the spectralline profiles during flares. MTs are generally found near the flareloop footpoints and persist only for a very short period of time.In this paper, we investigate WLFs and their relationshipwith MTs in NOAA active region 11515 by using data takenby the Helioseismic and Magnetic Imager ( HMI ; Scherrer et al.2012, Schou et al. 2012a, 2012b) and the
Atmospheric ImagingAssembly ( AIA ; Lemen et al. 2012) on board the
Solar DynamicsObservatory ( SDO ; Pesnell et al. 2012). We use HMI continuumintensity images to identify WLFs, and use the HMI line-of-sightmagnetograms to detect the magnetic field changes. Our obser-vations are presented in Section 2. Analysis and results are givenin Section 3. In Section 4 we present a brief summary and dis-cussion.
2. Observations
In this study, we use HMI full-disk continuum filtergrams andline-of-sight magnetograms, both observed by using the line ofFe I × I ′′ . AIA 131 Å , 171 Å and 1600 Åimages are used to show the morphology of the flares in the up-per atmosphere e.g., Figure 4). The AIA 1600 Å images are alsoused to determine the flare regions (see below). The temporalcadence of the AIA observations is 12 s in the extreme ultravio-let (EUV) passbands and 24 s in the ultraviolet (UV) passbands.The spatial resolution is about 1 . ′′ .It should be noted that the HMI continuum intensity ( I c ) isobtained by “reconstructing” the spectral line through the fol-lowing equation: I c = X j = [ I j + I d exp ( − ( λ − λ ) σ )] , (1)where λ , σ and I d are the rest wavelength, line width and linedepth, estimated by taking six sampling points ( I j ) across the Fe I absorption line 6173.3 Å (Couvidat et al., 2012).NOAA active region 11515 was very active, producingnearly one hundred flares during its passage on the solar disk.However, there were no X-class flares. In this study we only con-sider flares that occurred between E ◦ and W ◦ to avoid pos-sible problems with projection e ff ects. In total 70 flares above Cclass were recorded during this period.Usually it is easier to identify WL emission enhancement inflares with a high GOES class, whereas it is di ffi cult to detectWL emission enhancement in low-class flares since normallythe WL emission in such flares is only weakly enhanced abovethe quiescent values. To clearly show the small changes in theWL emission, we have constructed pseudo-intensity images bymagnifying the di ff erence between two adjacent continuum fil-tergrams by a factor of 5. The pseudo-intensity is expressed as: I ′ t + s = ( I t + s − I t ) × + I t , where I t and I t + s are the originalcontinuum intensity at two adjacent times with a gap of 45 s (theobserving cadence of HMI). From the pseudo-intensity images,it is much easier to detect WLFs even when they are very weak(Figure 1). To demonstrate the capability of our method, onlineanimations of a WLF using the original HMI continuum inten-sity images and using the pseudo intensity images, respectively,are provided. It should be noted that this method is only usedto help determine whether there is an impulsive enhancement ofWL emission during a flare.Through this approach, 20 out of the 70 flares are found toreveal an enhancement in the WL emission, though with dif-ferent sizes of the enhancement area. Following the definitionof WLFs from many recent investigations (e.g., Krucker et al.2015; Huang et al. 2016; Kuhar et al. 2016), we define these 20flares which show an impulsive enhancement in HMI continuumintensity as WLFs. Twelve of them are found in M-class flaresand the other eight are found in C-class flares. Table 1 lists theobservational information and results for these WLFs.
3. Analysis and results
The dates, peak times, GOES classes and positions of these 20WLFs can be seen from Table 1. For each WLF, we first definethe WLF region. To estimate the fluctuation of the background,
Num Date Peak GOES AR △ T wl S wl dI mwl dI awl dB l B l sign I wl Time Class Location (min) (1 ′′ ) (( I pwl − I wl ) / I wl ) (( I pwl − I wl ) / I wl ) (( B pl − B l ) / B l ) change ( I / I max )1 2012.07.03 17:02 C9.0 S17W08 3.75 10.50 9 . ± .
1% 6 . ± .
1% 1 .
4% Y 72 . . ± .
2% 8 . ± .
2% 1 .
7% Y 82 . . ± .
3% 5 . ± .
3% 5 .
7% Y 85 . . ± .
0% 7 . ± . − .
0% Y 66 . . ± .
3% 6 . ± . − .
0% N 84 . . ± .
2% 7 . ± . − .
7% N 50 . . ± .
3% 9 . ± . − .
2% Y 80 . . ± .
1% 7 . ± . − .
1% — 93 . . ± .
4% 6 . ± . − .
7% — 64 . . ± .
3% 6 . ± . − .
4% Y 74 . . ± .
1% 14 . ± . − .
2% Y 98 . . ± .
2% 6 . ± . − .
7% — 60 . . ± .
3% 10 . ± . − .
0% Y 91 . . ± .
1% 10 . ± . − .
1% Y 93 . . ± .
1% 7 . ± . − .
1% Y 91 . . ± .
3% 8 . ± .
3% 0 .
3% Y 90 . . ± .
2% 7 . ± . − .
3% Y 94 . . ± .
1% 10 . ± .
1% 13 .
5% Y 81 . . ± .
1% 7 . ± . − .
2% — 78 . . ± .
1% 6 . ± . − .
2% — 66 . △ T wl is the duration of the WLF. S wl is the area of the WLF region. I wl and B l are the intensity and line-of-sight magnetic field strength before the flare peak in the WLF region. I pwl and B pl are the intensityand line-of-sight magnetic field strength at the peak time of the flare in the WLF region. dI mwl refers to the maximum value of ( I pwl − I wl ) / I wl in the WLF region and dI awl refers to the average value. dB l is the average value of ( B pl − B l ) / B l in the WLF region. For the sign change of B l , ‘Y’ means clear sign change, ‘N’ means no sign change, and ‘—’ means not obvious. I is the average intensity of AIA 1600Å in the WLF region and I max is the maximum value in the whole flare region. we select three quiet-sun regions (R1, R2, R3; marked in Figure2(b) & (d)) outside the sunspots, which are located in the north,south and east of the active region, respectively. These three re-gions are nearly devoid of field during the periods when we es-timate the fluctuation. All three regions have a size of 20 × ff erence within each ofthese three regions during three di ff erent periods. From Figure 3we can see that the standard deviations in these three regions dur-ing di ff erent periods are very close. The average value is about0.013 and the maximum value is about 0.016, which can be re-garded as a level of the fluctuation of the background. We definethe WLF region as the area where the emission enhancement(( I pwl − I wl ) / I wl ) is greater than 0.05 ( > × . △ T wl ). Wecalculate dI awl , which is the average value of the percentage ofWL intensity increase at the peak of the flare relative to that be-fore the peak of flare. The maximum value of this percentage, dI mwl , is also calculated. The error of WL enhancement is esti-mated as the average of the standard deviations of the changes ofWL emission in the three quiet-sun regions (R1, R2, R3) at thesame time, and it reflects the intensity fluctuation of the gran-ules. Similar to dI awl , dB l is the average value of the percentagechange of unsigned line-of-sight magnetic field in the WLF re-gion. Also, we calculate the ratio between the average intensityof the WLF region in the AIA 1600 Å images, and the maximumintensity in the 1600 Å images within the whole flaring regionduring the peak time of the flare ( I wl ). Using this ratio, we candetermine whether the WLF occurred in the central area of theflare or elsewhere. The central area of a flare is defined as theregion where the AIA 1600 Å intensity is greater than half themaximum value at the peak time. It should be noted that the cen- tral area of a flare refers to the central area of the whole flaringregion.From Table 1 we can see that these WLFs generally have ashort duration, with an average lifetime of 4.65 minutes. The sizeof the observed WLFs are generally small, with an average areaof about 37.18 square arcseconds. From the percentage increaseof the continuum intensity ( dI awl , dI mwl ), we can see that the WLenhancements in these WLFs are generally very weak. Hudsonet al. (2006) detected 11 WLFs in the WL channel of TRACE,including 4 C-class flares. The minimum excess contrast is only0.08 ± RHESSI and HMI, Kuhar et al. (2016) studied 43 WLFs(M- & X- classes) and found that the lowest change of white-light emission is 0.08 ± dI mwl in Table 1, is greater than 20% for only 5 WLFs in our study.The lowest value of dI mwl is 8% ± . dI awl ) are mostly less than 10%, with an averageof 8.1% and a minimal of 5 . ± . . ± . I wl for all WLFs are greaterthan 0.5, meaning that the WL enhancements all occur in thecentral areas of the flare ribbons defined by the enhanced 1600Å emission.Figure 2 shows the occurrence times and locations of theseWLFs. In panel (a) the solid line is the soft X-ray (1-8Å) fluxmeasured by GOES. The dotted vertical lines mark the peaktimes of the 70 flares detected in AR 11515 between E ◦ ∼ W ◦ on the solar disk, where the red and black ones are for theWLFs and normal flares, respectively. Panels (b) and (d) are thecontinuum intensity images at two di ff erent times. Panels (c) and(e) are the corresponding line-of-sight magnetgrams. In panels (b)-(e) the red circles mark the central positions of these WLFs.From panel (a) and Table 1 we can see that these WLFs occurmainly on the days of July 4th and July 5th. From panels (b)-(e)we can see a narrow ribbon-like magnetic field structure in activeregion 11515 on these two days. The magnetic field at the ribbonis negtive, but the magnetic field on both sides of the ribbon arepositive. All the WLFs are distributed along this ribbon.Figure 4 shows the white-light di ff erence images for these 20WLFs. The red contours in each panel mark the WLF region wedefined based on the method described above. All these imageshave a size of ∼ ′′ × ′′ , corresponding to the main flaring re-gion and covering the whole WL enhancement regions. It shouldbe noted that the WL di ff erence images are obtained by calcu-lating ( I pwl − I wl ) / I wl , where I pwl is the HMI continuum intensity atthe peak time of the flare and I wl is the intensity several minutesbefore the flare peak.Figure 5 shows an individual example for an M4.7 WLFwhich occurred on July 5th. Panels (a)-(c) show the AIA 1600Å,171Å and 131Å images at the peak time of this WLF, respec-tively. From these three panels we can see that the spatial scaleof this flare is small. The white box in panel (a) corresponds tothe field of view (FOV) in panels (d)-(i). Panels (d) and (e) arethe continuum images at the time before and at the peak time ofthe flare. Panel (f) is the di ff erence image between them. Panels(g)-(i) are similar to panel (d)-(f) but for the magnetic field. Thegreen box in panels (d)-(i) shows the region where the WLF andMT occurred. From these panels we can see that there is a sig-nificant enhancement of the continuum intensity in a very shorttime period and the HMI line-of-sight magnetic field changessignificantly at the same time and the same locations. In otherwords, WLF and MT occurred simultaneously and co-spatially.Note that the changes of both the continuum intensity and mag-netic field are transient and not permanent.Figure 6 shows the temporal evolution of the white lightintensity (red) and unsigned line-of-sight magnetic field (blue)around the flaring times of these 20 WLFs. Each diamond onthese curves represents the average value of continuum intensityor unsigned magnetic field in the WLF region. The green shadedregion in each panel marks the duration of the WLF, which isdefined as the period between times when the WL enhancementobviously appeared and vanished. To estimate the uncertainty ofthe WL intensity, we first plot the WL curves in three quiet-sunregions (R1, R2, R3) during the same period of the flare. Thenwe take the average value of the three standard deviations de-rived from the three WL curves as the intensity uncertainty forthe WLF. The figure clearly shows that the MTs and WLFs oc-cur simultaneously. From Table 1 we can see that MTs in 13WLFs show an obvious change of sign. It should be noted thatthe MT in the WLF 12 (the C9.1 WLF on Jul 5th) seems to bevery weak and di ffi cult to be identified. This is because the tran-sient mainly occurred in regions of negative field, the strengthof which is much weaker compared to that of the positive field.Thus, the MT is not reflected obviously in the curve of unsignedmagnetic field for this WLF. Also, we find a close relationshipbetween the WL enhancement and magnetic field change, whichis consistent with the conclusion of Song & Zhang (2016) thatthe magnetic field and intensity variations are closely related andthey are possibly the two facets of the same phenomena of a solarflare. It should be noted that the WL enhancement and magneticfield change in their observations are permanent. MTs are usu-ally believed to be artifacts due to the distortions of line profile(Patterson 1984, Ding et al. 2002, Qiu & Gary 2003, Isobe et al.2007, Mauraya et al. 2012), although some recent studies sug- gest that some of them are real (Matthews et al. 2011, Harker &Pevtsov 2013). In this study, we do not investigate whether MTsare real or not, which relies on a detailed future examination ofthe profile of the 6173 Å line.In Figure 7, panel (a) shows the relationship between thechange of WL emission ( dI awl ) and the AIA 1600Å intensity( I ) in WLF regions, while panel (b) shows the relationshipbetween the change of WL emission ( dI awl ) and the absolutechange of unsigned line-of-sight magnetic field ( | dB l | ). In panel(a), we can see that the values of I wl in all WLFs are greaterthan 0.5, which means that all of these WLFs occur at the centralregions of the flare ribbons. There is a trend that the location ofthe WLF is closer to the center of the whole flaring region whenthe enhancement of the WL emission is stronger. Panel (b) re-veals a possible linear correlation between the changes of WLemission and unsigned magnetic field. A greater enhancementof the WL emission is often accompanied by a larger change ofthe line-of-sight magnetic field. This can be seen from Table 1.
4. Summary and discussion
WLFs were believed to be very rare compared to the frequentoccurrence of solar flares. However, in this study we find atleast 20 WLFs out of the 70 flares in NOAA AR 11515 duringits passage on the solar disk ( E ◦ ∼ W ◦ ). Thus, the occur-rence rate of WLFs in this active region is at least 28.6%, whichprovides further evidence towards the idea that all flares mayhave a WL component, a possibility discussed by Hudson et al.(2006) several years ago. Our findings also emphasize the factthat WLFs can occur in less energetic flares, i.e, C- and lowerM-class flares, following some previous studies (e.g., Matthewset al. 2003; Hudson et al. 2006; Jess et al. 2008; Kowalski et al.2015a). For most of these 20 WLFs, the WL enhancement lastsfor a short duration and occurs in a small region. The averageenhancement ( dI awl ) of the WL emission in the WLF region isgenerally very small, with an average of 8.1%. If we regard aWLF whose change of WL emission ( dI awl ) is lower than 10% asa weak one, then there are 17 weak WLFs in our sample and thepercentage is 85% (17 / I wl > .
5) ofthe flare ribbons. This can be understood as the central areas offlare ribbons usually have a greater energy deposition and WLenhancement most likely occurs there. All 20 WLFs in NOAAAR 11515 occur mainly on 4 July and 5 July 2012, and are dis-tributed along a narrow ribbon with negative magnetic flux sur-rounded by positive flux on both sides. The high occurrence rateof WLFs in this active region may be related to the developmentof this special magnetic field configuration. Our results appearto support the conclusion of Neidig & Cliver (1983) that WLFsare often produced in large and magnetically complex active re-gions.Another interesting result is that these 20 WLFs are all ac-companied by MTs. For the remaining 50 flares without obvi-ous WL enhancement, there are no detectable MTs. It is stillan open question whether MTs are real or not (see Harker &Pevtsov (2013) and references therein). Nevertheless, our obser-vations suggest that MTs and WLFs are closely related when weconsider the corresponding HMI data for this particular activeregion. They occur at the same time and at the same location.A larger enhancement of WL emission is often accompanied by a lager change of the line-of-sight component of the unsignedmagnetic field. Considering the close relationship between MTsand WLFs, many previously reported flares with MTs may beWLFs.Although many more WLFs are detected in AR 11515,they are mostly small-scale and short-lived brightenings in the6173 Å continuum, similar to the small WL kernel reported byJess et al. (2008). In the future we plan to examine more ARsto see if most WLFs possess similar features. Also, it should benoted that the HMI continuum intensity ( I c ) is obtained by “re-constructing” a spectral line. Although the enhancement in theHMI continuum intensity is commonly identified as a signatureof WLFs (e.g., Krucker et al. 2015; Huang et al. 2016; Kuhar etal. 2016), we cannot exclude the possibility that the HMI I c valuemay have been, to a certain degree, a ff ected by the emission inthe line core.In addition, several questions need to be clarified with fu-ture observations. First, what is the energy source of such small-scale and short-lived WLFs? If they are powered by electronsbeams, as in the type-I WLFs studied previously, they shouldaccompany small bursts in the hard X-ray (or microwave) emis-sion. If the heating energy is from other sources like Alfv´enwaves (Fletcher & Hudson 2008) or even in situ energy release,there may not be obvious correlation between WL enhancementand the hard X-ray emission. The latter case belongs to type-II WLFs. Simultaneous observations of spectral lines formedat di ff erent layers are helpful in judging the energy source. Toanswer this question coordinated observations of several instru-ments with very high cadence and spatial resolution are required.Second, what is the relationship between such a small WL ker-nel and the whole flare? What special condition is required toproduce a WLF? For this purpose, vector magnetic field obser-vations and three-dimensional magnetic field extrapolation arerequired to check if the WLFs are located at peculiar sites whereeither magnetic reconnection likely occurs or energy can be eas-ily transported there. Acknowledgements.
This work is supported by NSFC grants11790304(11790300), 11125314 and 11733003, the Recruitment Programof Global Experts of China, the Specialized Research Fund for State KeyLaboratories and the Max Planck Partner Group program. H.T. acknowledgessupport of ISSI and ISSI-BJ to the team “Diagnosing heating mechanisms insolar flares through spectroscopic observations of flare ribbons”. We thank theanonymous reviewer for the carefully reading and very constructive comments.
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HMI continuum intensity images of an M5.3 WLF observed at 09:54:53 UT on 2012 July 4. (a) is the original intensity and(b) shows the pseudo-intensity. Red arrows point out the regions where the WL enhancement occurred. (An animation of this figureis available.)
Fig. 2.
White-light flares (WLFs) in NOAA AR 11515. ( a ) 70 flares (C- and M- classes) occurred between E ◦ to W ◦ . The solidline is GOES soft X-ray (1-8 Å). The red dashed lines mark the peak times of WLFs, and the black dotted lines mark these of normalflares. ( b )-( e ) Locations of these WLFs in the active region. The circles in panels (b) and (c) show WLFs 1–12 while the circlesin panels (d) and (e) show WLFs 13–20. ( b ) and ( d ): HMI continuum images at two di ff erent times, ( c ) and ( e ): HMI line-of-sightmagnetgrams at two di ff erent times. The black boxes R1, R2 and R3 in panels (b) and (d) mark the quiet-sun regions selected toestimate the error in the measured continuum intensity. The size of each box is 20 ×
20 pixels.
Fig. 3.
Standard deviations of the di ff erence of the HMI continuum intensity in three quiet-sun regions (R1, R2 and R3) duringthree di ff erent periods (top, middle, bottom). The red, blue and green colors represent the results in the regions of R1, R2 and R3,respectively. Fig. 4.
Snapshots of 20 WLFs. The images are obtained by calculating ( I pwl − I wl ) / I wl , where I pwl and I wl are the intensity imagestaken at two di ff erent times. The red contours mark the WLF regions. Each image has a size of ∼ ′′ × ′′ . Fig. 5.
A white-light flare in NOAA AR 11515. ( a ), ( b ) and ( c ) are AIA 1600, 171 and 131 Å images at the peak time of this WLF,respectively. ( d ) and ( e ) are HMI continuum images at the beginning and peak time of this WLF, respectively. ( f ) is the di ff erenceimage between ( d ) and ( e ). (g) and (h) are HMI line-of-sight magnetgrams at the beginning and peak time of this WLF, respectively.( i ) is the di ff erence image between ( g ) and ( h ). The white box in panel (a) corresponds to the FOV in panels (d)-(i). The green boxin panels ( d )-( i ) marks the location where the flare occurred. Fig. 6.
Temporal evolution of the HMI continuum (red) and unsigned B l (line-of-sight magnetic field, blue) for the 20 WLFs. Theblack dotted line marks the time of the WL emission peak for each WLF. The green shaded region marks the duration of each WLF.The time range shown in each panel is from 25 minutes before the peak time to 25 minutes after it. Fig. 7. (a) The relationship between WL enhancement ( dI awl , the average enhancement of WL emission in the WLF region) andAIA 1600 Å intensity ( I wl , the ratio between the average AIA 1600 Å intensity in the WLF region and the maximum 1600 Åintensity in the whole flare region) for 20 WLFs in NOAA AR 11515. The positive correlation indicates that a WLF has a strongerWL emission if it occurs closer to the center of the whole flaring region. (b) The relationship between WL enhancement ( dI awl ) andmagnetic field change ( | dB l | , the absolute change of the unsigned B l in WLF region). A greater enhancement of the WL emission isoften accompanied by a larger change of the line-of-sight magnetic field.in WLF region). A greater enhancement of the WL emission isoften accompanied by a larger change of the line-of-sight magnetic field.