Spectral and Imaging Observations of a C2.3 White-Light Flare from the Advanced Space-Based Solar Observatory (ASO-S) and the Chinese H$\alpha$ Solar Explorer (CHASE)

Solar white-light flares are characterized by an enhancement in the optical continuum, which are usually large flares (say X- and M-class flares). Here we report a small C2.3 white-light flare (SOL2022-12-20T04:10) observed by the \emph{Advanced Space-based Solar Observatory} and the \emph{Chinese H$\alpha$ Solar Explorer}. This flare exhibits an increase of $\approx$6.4\% in the photospheric Fe \textsc{i} line at 6569.2\,\AA\ and {$\approx$3.2\%} in the nearby continuum. The continuum at 3600\,\AA\ also shows an enhancement of $\approx$4.7\%. The white-light brightening kernels are mainly located at the flare ribbons and co-spatial with nonthermal hard X-ray sources, which implies that the enhanced white-light emissions are related to nonthermal electron-beam heating. At the brightening kernels, the Fe \textsc{i} line displays an absorption profile that has a good Gaussian shape, with a redshift up to $\approx$1.7 km s$^{-1}$, while the H$\alpha$ line shows an emission profile though having a central reversal. The H$\alpha$ line profile also shows a red or blue asymmetry caused by plasma flows with a velocity of several to tens of km s$^{-1}$. It is interesting to find that the H$\alpha$ asymmetry is opposite at the conjugate footpoints. It is also found that the CHASE continuum increase seems to be related to the change of photospheric magnetic field. Our study provides comprehensive characteristics of a small white-light flare that help understand the energy release process of white-light flares.

Spectral and Imaging Observations of a C2.3 White-Light Flare from the Advanced Space-Based Solar Observatory (ASO-S) and the Chinese Hα Solar Explorer (CHASE) Abstract Solar white-light flares are characterized by an enhancement in the optical continuum, which are usually large flares (say X-and M-class flares).
Here we report a small C2.3 white-light flare (SOL2022-12-20T04:10) observed by the Advanced Space-based Solar Observatory and the Chinese Hα Solar Explorer.This flare exhibits an increase of ≈6.4% in the photospheric Fe i line at 6569.2 Å and ≈3.2% in the nearby continuum.The continuum at 3600 Å also shows an enhancement of ≈4.7%.The white-light brightening kernels are mainly located at the flare ribbons and co-spatial with nonthermal hard Xray sources, which implies that the enhanced white-light emissions are related

Introduction
Solar white-light flares (WLFs) are identified by a sudden increase in the visible continuum (Neidig and Cliver, 1983;Neidig, 1989).The first WLF (also the first solar flare) was reported by Carrington (1859) and Hodgson (1859).Most WLFs occur in the vicinity of sunspots and manifest as brightening kernels, but some can also occur in a region almost without sunspot (e.g.Hudson et al., 1994) and appear as loop-like structures (e.g.Hudson et al. 1992;Hudson, Wolfson, and Metcalf 2006;Jejčič, Kleint, and Heinzel 2018).While it is generally accepted that WLFs constitute a minority in the overall flare family (Fang et al., 2013) with X-and M-class flares (i.e.major flares) being the most energetic, there are different views suggesting that all flares may exhibit white-light brightenings (e.g.Matthews et al., 2003;Jess et al., 2008;Song et al., 2018).This hypothesis gains support from the discovery of white-light continuum enhancements in small flares.
Up to now, there have been only about 20 C-class WLFs reported including several low C-class ones from both statistical studies (e.g.Matthews et al., 2003;Hudson, Wolfson, and Metcalf, 2006;Song et al., 2018;Song and Tian, 2018;Castellanos Durán and Kleint, 2020) and case analyses (e.g.Jess et al., 2008;Song et al., 2020).Statistically, C-class WLFs exhibit an average increase of ≈10% in the visible continuum.Some C-class WLFs with a lower magnitude could display a relatively higher white-light enhancement (e.g.Hudson, Wolfson, and Metcalf, 2006;Song et al., 2018;Song and Tian, 2018), which seems to deviate from the behavior as observed in major flares.Note that this discrepancy may be attributed to sample bias and instrumental limitation.For case studies, Jess et al. (2008) observed a C2.0 WLF showing a white-light kernel with a diameter of ≈0.4 ′′ using the Swedish Solar Telescope (SST) among the highest resolution telescopes.This small scale WLF demonstrates an unusually high increase of ≈300% in the blue continuum, surpassing a typical enhancement in WLFs.However, degrading the spatial resolution to 1 ′′ decreases the enhancement to only 1%.In another case study, Song et al. (2020) presented a C2.3 WLF with an enhancement of ≈18% observed by the Helioseismic and Magnetic Imager (HMI: Scherrer et al., 2012) on the Solar Dynamics Observatory (SDO: Pesnell, Thompson, and Chamberlin, 2012).This particular flare is associated with weak hard X-ray (HXR) emissions, with the white-light enhancement unlikely caused by accelerated electron beams.Its magnetic field morphology and evolution are more supportive of magnetic reconnection in the lower atmosphere (e.g.Ding, Fang, and Yun, 1999;Chen, Fang, and Ding, 2001) rather than in the corona.This suggests that both the non-thermal particle injection and evolution of magnetic fields need to be considered in studying WLFs.Given the scarcity of studies on small WLFs, more attention should be paid to the WLFs with a low magnitude.
Previous studies have indicated that permanent changes in the photospheric magnetic field commonly appear in flares (e.g.Sudol and Harvey, 2005;Castellanos Durán, Kleint, and Calvo-Mozo, 2018).These changes mean that the magnetic components undergo abrupt transformations and do not return to their original levels for a long time.The coronal implosion process (Hudson, 2000) predicts that photospheric fields become more horizontal during flares, thereby changing the line-of-sight (LoS) magnetic field strength [B LoS ] (Hudson, Fisher, and Welsch, 2008).This has been used for an interpretation of the appearance of B LoS changes [∆B LoS ] (e.g.Petrie and Sudol, 2010;Gosain, 2012;Sun et al., 2012).However, differences in ∆B LoS between the chromosphere and photosphere suggested that the chromospheric ∆B LoS may not support a contraction of the field line in flare loops (e.g.Kleint, 2017).Recently, a statistical study (Castellanos Durán and Kleint, 2020) showed that white-light brightenings and ∆B LoS in the photosphere can overlap less than 60%, and the areas of these two parameters appear to have a power-law relation.The relationship between ∆B LoS and white-light brightenings deserves a further study, especially for small WLFs.
SOLA: ms.tex; 3 May 2024; 1:16; p. 3 Spectral analysis is a good way to study the dynamics of plasma in WLFs, especially via some Fe i lines formed in the photosphere (e.g.Babin and Koval, 1992;Lin, Zhang, and Zhang, 1996;Hong et al., 2018).Jurčák et al. (2018) reported a WLF with the response of Fe i lines at 6301.5 Å and 6302.5 Å, which was estimated to originate in the chromosphere instead of photosphere.In addition, the Hα line at 6562.8 Å formed from the photosphere to chromosphere can help us understand the physical process of WLFs.This line mainly behaves as an emission profile with a central reversal and a red asymmetry during WLFs (e.g.Babin and Koval, 1992;Zhou, Fang, and Hu, 1997).
In this work, we report a small C2.3 WLF (SOL2022-12-20T04:10) observed by the Advanced Space-based Solar Observatory (ASO-S: Gan et al., 2023) and the Chinese Hα Solar Explorer (CHASE: Li et al., 2022).ASO-S provides images in the white-light continuum at 3600 Å and also HXR spectra at ≈10 -300 keV.CHASE has spectral observations in the Fe i 6569.2Å and Hα 6562.8Å lines.All these allow us to study the comprehensive characteristics of this C2.3WLF when combining with some other multi-wavelength observations.In the following Section 2, we describe the data reduction.Our results are presented in Section 3, followed by a summary and discussion in Section 4.

Data
The data used in this work are obtained by multiple telescopes.The main telescopes include the Lyman-alpha Solar Telescope (LST: Li et al., 2019;Chen et al., 2019;Feng et al., 2019) aboard ASO-S, CHASE, and HMI on SDO.LST consists of three instruments, one of which is the White-light Solar Telescope (WST).WST can provide full-disk images in the continuum at 3600±20 Å.The pixel size of the images is ≈0.5 ′′ while the spatial resolution is about 4 ′′ (Jing et al., 2024).The cadence is two minutes in a routine mode.The level-2.5 (L2.5) data of WST are used here, which have been done for scientific calibration and image registration.CHASE provides spectral data of the Fe i line at 6569.2 Å (6567.8-6570.6Å) and the Hα line at 6562.8 Å (6559.7 -6565.9Å) via a raster scanning mode (RSM).The data have a cadence of ≈1 minute, a pixel size of ≈1 ′′ , and a spectral resolution of 0.048 Å pixel −1 , with a two-binning mode in both space and spectrum.The calibrated L1.5 data are used here (Qiu et al., 2022).The full-disk images of Fe i 6173 Å and the LoS magnetogram presented in the work are obtained from SDO/HMI, which have a pixel size of 0.5 ′′ and a cadence of 45 seconds.
The Hard X-ray Imager (HXI: Zhang et al., 2019;Su et al., 2022) aboard ASO-S is an HXR imaging spectrometer observing the Sun in an energy range of ≈10 -300 keV with energy and angular resolutions of ≈16.5% at 32 keV and 3.1 ′′ , respectively.The HXI L1 data are used in the present study.The Spectrometer Telescope for Imaging X-rays (STIX: Krucker et al., 2020) on Solar Orbiter (SolO: Müller et al., 2020) has energy and angular resolutions of ≈1 keV @6 keV and ≥7 ′′ , respectively, in a range from 4 keV to 150 keV.Its L1 pixel and SOLA: ms.tex; 3 May 2024; 1:16; p. 4 spectrogram data are used here.At the observation time of the flare event, STIX had a separation angle of 19.2 • with Earth and a distance of 0.9 AU to the Sun.The Atmospheric Imaging Assembly (AIA: Lemen et al., 2012) on SDO provides the images for flare loops in the Extreme ultraviolet (EUV) channel at 131 Å with a pixel size of 0.6 ′′ and a cadence of 12 seconds.The SXR 1 -8 Å light curve is from the X-Ray Sensor (XRS: Hanser and Sellers, 1996) aboard the Geostationary Operational Environmental Satellite (GOES) with a high cadence of 1 second.

Data Reduction
We make a co-alignment for the images from different instruments via the routines in the SolarSoftWare (SSW: Freeland and Handy 1998).Firstly, the AIA and HMI images are registered using aia_prep.pro.The solar rotation has been removed via drot_map.pro.Then the white-light images from WST and CHASE are co-aligned with the HMI continuum images using xyoff.probased on cross correlation for sunspot features.The uncertainty of the co-alignment is estimated to be about 1 ′′ .
For HXI light curves at 20 -50 keV and 50 -100 keV, we use the combined data of three total flux detectors (D92, D93, and D94) with a cadence of four seconds.The images are reconstructed using HXI CLEAN algorithm (Su et al., 2019b) with sub-collimators G3 -G10, which correspond to a spatial resolution of ≈6.5 ′′ .We also use STIX data to produce the light curves at 4 -10 keV and 10 -20 keV with a cadence of four seconds and to reconstruct images via the Expectation Maximization (EM) algorithm (Massa et al., 2019).In addition, we make spectral fitting from STIX via a thermal component plus a nonthermal thick target component.Note that the light-travel time at Earth with a correction of 35 s for STIX has been considered.
In order to remove non-flaring contributions and calculate the relative enhancements caused by flare, we subtract a background or base image from flaring images for different wave bands.For HMI and WST images, an average for 30 minutes before the flare onset is adopted as the background.For CHASE images, the last moment at 04:19:48 UT, i.e. around the flare end time, is selected as the background which is relatively quiet during the whole observation of CHASE.Note that this may lead to an underestimation of the CHASE continuum enhancement to some extent.When estimating the uncertainty of relative enhancements for each wave-band emission, we choose a quiet-Sun region with a size of 100 ′′ ×16 ′′ (X = [−799 ′′ , −699 ′′ ], Y = [483 ′′ , 499 ′′ ]) to calculate its relative intensity fluctuations, i.e. the error bars as plotted in Figures 3a and b.Note that this quiet-Sun region is out of the field of view of the images as shown in Figures 1d -i.
Using the CHASE spectral data, we can measure the Doppler velocities of Fe i and Hα lines and also calculate the Hα asymmetry.The reference line centers of Fe i and Hα are determined by averaging the profiles over the quiet-Sun region as mentioned above.The excess profile, i.e. after subtracting the background profile, is adopted to obtain the Doppler velocity.For Fe i, a single Gaussian fit with a linear background is used and the average of the background is further SOLA: ms.tex; 3 May 2024; 1:16; p. 5 used to define the nearby continuum.It is found that the intensity of this defined continuum is approximately equal to that at the wavelength of 6568.6 Å.For Hα, bisector (e.g.Chae et al., 2013) and moment analyses are applied to derive its Doppler velocity.The velocity uncertainties of Fe i and Hα lines are estimated to be within ±0.6 and ±3.9 km s −1 , respectively.The Hα asymmetry is calculated based on its original profiles.Referring to Equation (1) in Asai et al. (2012), the asymmetry is expressed as follows, where I rp and I bp denote the peaks of Hα red and blue wings, respectively.A positive/negative RA represents a red/blue asymmetry.The uncertainty of the asymmetry is estimated to be about ±0.014.When showing the profiles of Fe i and Hα (in Figures 3c -f and 4a -d), we make a normalization by using the nearby quiet-Sun continuum, i.e. in a contrast way.More specifically, the nearby quiet-Sun continuum intensity of the Hα line is determined by averaging the intensities of the ten wavelength points in its bluest wing, while that of the Fe i line is defined as the average of the linear background of this line.
To investigate the relationships of white-light emission with B LoS and ∆B LoS in the photosphere, the magnetograms from HMI are reshaped into the same pixel size as CHASE images.Then a stepwise function (Sudol and Harvey, 2005) describes the temporal evolution of B LoS as expressed by in which t represents time and a, b, c, n, and t 0 are free parameters to fit.From this stepwise function, the characterized start time [t s = t 0 − πn −1 ] and end time [t e = t 0 + πn −1 ] can also be obtained.∆B LoS is derived mainly via the method in Castellanos Durán, Kleint, and Calvo-Mozo (2018), which could be expressed as: where t max and t min represent the times of the maximum and minimum of B LoS , respectively, and the function sgn(x) represents the positive or negative sign of x.Here we also modify some criteria as follows.
1.In this short-duration event, the time range of fit is reduced to 60 minutes centered at the GOES peak time.2. The duration parameter [πn −1 ] must be smaller than the fit period, otherwise ∆B LoS = 0. 3. t s must be earlier than t e , otherwise ∆B LoS = 0. 4. Both t s and t e must be within the fit period, otherwise ∆B LoS = 0.
Note that the noise level of B LoS and ∆B LoS is determined as about ten Gauss (Liu et al., 2012).

Overview of the C2.3 Event
The event under study is a GOES C2.3 flare (SOL2022-12-20T04:10) occurring in NOAA active region 13171 (N26E59).This flare began at 03:46 UT, peaked at 04:10 UT, and ended at 04:19 UT on 20 December 2022 (see the GOES SXR 1 -8 Å light curve as shown in Figure 1a), showing a relatively long slow rise before 04:04 UT.CHASE observations covered the main period of the flare from 04:04 to 04:20 UT (denoted by two green vertical lines in Figure 1a), which we focus on in the following.Figure 1b plots the HXR light curves at different energy bands from STIX and HXI, together with the SXR 1 -8 Å flux and its temporal derivative.We can see that the HXR 10 -20 as well as 20 -50 keV emissions peak around the same time as the SXR temporal derivative, indicating the Neupert effect (Neupert, 1968).Note that the flare has no response in HXR emission above ≈50 keV. Figure 1c shows multiple emission curves integrated over the flaring core region as marked by the green box in Figures 1d -i.It is seen that the emissions at WST 3600 Å and CHASE Hα exhibit a notable rise followed by a decay during the flare.Note that their emission peaks may not represent the true peaks due to a low cadence of more than one minute.It is interesting that the CHASE and HMI continuum emissions show a similar trend, though having multiple peaks during the observation period.It should be mentioned that the main peaks of CHASE and HMI emissions appear in the main phase of the flare.
Figures 1d -i display multi-wavelength images of the flaring region around the flare peak time.We find two main white-light brightening kernels, referred to as K1 and K2, from the base-difference images of HMI, CHASE, and WST continua (Figures 1d -f).These two kernels are located near a small pore (indicated by a black arrow) as seen in the HMI and CHASE continuum images in Figures 1d  and e.From the HMI magnetogram in Figure 1g, we can see that K1 is located in negative polarities with a relatively strong magnetic field, while K2 seems to lie in a mix of positive and negative polarities with a weak field, which may be influenced by a projection effect.These two kernels are co-spatial with the Hα ribbons as shown in Figure 1h.K1 is located in the center of the south ribbon with a bright Hα source, while K2 lies on the edge of the north ribbon with a relatively weak Hα brightening.From the AIA 131 Å image in Figure 1i, one can see that K1 and K2 connect the same set of flare loops, which are likely to be a pair of conjugate footpoints.

Spatial Relationship of the White-light Kernels with Nonthermal HXR Sources
Figure 2 shows the HXR spectral imaging and fitting from STIX and HXI for the C2.3 flare.From the spectral fitting (Figure 2e), we can see a nonthermal component above ≈20 keV.Thus, we reconstruct the HXR images at 20 -35 keV from HXI for the main nonthermal sources.In addition, we reconstruct the HXR images at 16 -28 keV and 4 -10 keV from STIX for nonthermal and also thermal sources.These HXR sources corresponding to the flare ribbons and also  2a and  b), the nonthermal HXR sources at 20 -35 keV are mainly cospatial with the flare ribbons.In particular, the two white-light kernels, K1 and K2, match the nonthermal sources well.It is noted that K1 is close to the nonthermal source with a larger HXR flux while K2 corresponds to the one with a lower flux.At a later time of 04:09:38 UT (Figures 2c and d), the HXR sources evolve a little while K1 is still close to one of the nonthermal HXR sources.These results suggest that the white-light emissions are related to a nonthermal electron-beam heating although the timings of white-light emissions and HXR sources are not conclusive due to a low cadence of the white-light data from WST and CHASE.

Spectral Features of the White-light Kernels
Figures 3a and b show the temporal evolution of the relative enhancements of continuum and spectral line emissions from HMI, WST, and CHASE, which are integrated over an area of 3 ′′ ×3 ′′ for the two white-light kernels K1 and K2.One can see that at K1 (Figure 3a), all the enhancement curves exhibit an evident rise and then a decay during the flare.The maximum enhancements or increases of HMI, CHASE, and WST continua at ≈04:10 UT are 3.9%, 3.2%, and 4.7%, respectively.For comparison, the CHASE Fe i line core and integrated Hα line over ±2.3 Å have larger maximum increases of 6.4% and 38.0%, respectively.At K2 (Figure 3b), the maximum increases of HMI, CHASE, and WST continua, as well as CHASE Fe i core and Hα line for their main peaks at ≈04:10 UT are 4.3%, 3.0%, 1.9%, 2.2%, and 20.2%, respectively, most of which are smaller than the corresponding values at K1.Note that the other peaks in the HMI and CHASE continua, say, before 04:08 UT and after 04:15 UT, might be unrelated to the flare itself, as no obvious response is found in the HXR emissions.Figures 3c -f exhibit the temporal evolution of CHASE Fe i and Hα line profiles averaged over an area of 3 ′′ ×3 ′′ at K1 and K2.One can see that at both kernels, the Fe i line remains an absorption profile with a symmetric Gaussian shape during the flare (Figures 3c -d).Its line core as well as the nearby continuum have an increase followed by a decrease in the intensity over time.The Hα line varies from an absorption to an emission and then back to an absorption profile (Figures 3e -f).Its emission profiles display a central reversal with double peaks in line wings, which are not symmetric.K1 shows a red asymmetry but K2 has a blue asymmetry, namely the conjugate footpoints have an opposite asymmetry in Hα.
We further select the Fe i and Hα line profiles with maximum intensities at 04:10:19 UT to derive the Doppler velocity as shown in Figures 4a -d.Note that the excess profiles are adopted to do a Gaussian fit for Fe i and to make a moment plus a bisector analysis for Hα.One can see that the Fe i line at K1 exhibits a trivial redshift velocity which is actually within the uncertainty (Figure 4a).By contrast, the excess Fe i emission at K2 has an evident redshift with a velocity of 1.70 km s −1 (Figure 4b), indicative of a plasma downflow at the lower atmospheric layer.The excess Hα profiles with a red or blue asymmetry at K1 and K2 show a much larger redshift or blueshift velocity (Figures 4c -d).
Moreover, we derive the Hα Doppler velocity map via moment for the time at 04:10:19 UT as shown in Figure 4e.To get a reliable result, we artificially set the velocity to 0 when the peak intensity of the excess profile in a pixel is less than 100 DNs.It is seen that evident Doppler velocities mainly come from the two brightened flare ribbons as marked by green contours (same as the white contours in Figure 1h), whose excess intensities are larger than 500 DNs.In particular, Hα can show redshift and blueshift velocities at the conjugate footpoints, which has been revealed from the profiles at K1 and K2 as described above.We also plot the asymmetry map in Figure 4f.One can see that only at the flare ribbons, the Hα profiles show a notable red or blue asymmetry with an obvious line-wing enhancement.Note that the asymmetry distribution looks quite similar to that of Doppler velocity at the two ribbons, which can also be revealed by the scatter plot in Figure 6i.However, one should caution that the red (blue) asymmetry of Hα profiles, or the measured redshift (blueshift) velocity is not necessarily caused by a plasma downflow (upflow), as already revealed in previous studies (e.g.Gan, Rieger, and Fang, 1993;Heinzel et al., 1994;Ding and Fang, 1996;Kuridze et al., 2015).It is also interesting to see that the asymmetry is reversed for the conjugate footpoints at K1 and K2.As given in Figures 4c and d, the asymmetries of Hα profiles at K1 and K2 are 0.079 and −0.036, respectively.

Photospheric Magnetic Field Changes at the White-light Kernels
Figure 5a shows the map of photospheric magnetic field change, i.e., ∆B LoS , for the flare region.We can see some notable changes in B LoS (high ∆B LoS ) near/at the white-light kernels K1 and K2.Note that some prominent changes of B LoS also appear in some other regions without white-light brightenings, which is actually similar to the study of Castellanos Durán and Kleint (2020).We further plot the temporal evolution of B LoS at K1 and K2 as well as the fit using a stepwise function in Figures 5b and c.It is seen that there seems to be no obvious change in B LoS , i.e. with a poor fit, at K1, while at K2, there exists an evident change in B LoS with a value of 40.5 Gauss during the flare.In fact, the relationship between the white-light emission and B LoS in this small WLF is somewhat similar to the results in Castellanos Durán and Kleint (2020).

Statistical Relationships among Different Parameters from Flare Ribbons
We check the statistical relationships among different parameters within the two flare ribbons (marked by the contours in Figures 1h, 4e and f, and 5a).
To increase the reliability, here we collect all the available time frames in observations, namely four frames for the CHASE continuum around its maximum, WST continuum, HMI B LoS , and CHASE Hα asymmetry but only one frame for

Summary and Discussions
In this paper, we perform a comprehensive study of a small C2.3 WLF observed by CHASE and ASO-S, combining with the multi-band observations from SDO and SolO.Our main results are summarized as follows.
1.This small flare shows relative enhancements or increases of 4.3%, 3.2%, 4.7%, and 6.4% in the HMI, CHASE, and WST continua and CHASE Fe i line, respectively.The white-light brightening kernels are mainly located near a small pore as well as within the Hα flare ribbons.2. At the white-light kernels, the CHASE Fe i line remains an absorption profile with a symmetric Gaussian shape during the flare.It can show a redshift velocity up to 1.7 km s −1 .By contrast, the Hα line changes from an absorption to an emission profile, the latter of which has a central reversal with a linewing enhancement.In addition, the Hα emission profile displays a red or blue asymmetry corresponding to a velocity of several to tens of km s −1 .3. Due to a low cadence of HMI, CHASE, and WST observations, the peak times of white-light emissions and their temporal relationship with the HXR emission are inconclusive in this flare.However, the white-light kernels are found to be co-spatial with the nonthermal HXR sources at 20 -35 keV.This suggests that the white-light emissions are related to a nonthermal electronbeam heating, whose radiative backwarming effect could not be excluded.4. One of the identified white-light kernels is accompanied by a change of photospheric magnetic field [∆B LoS ], while the other is not.Only the CHASE continuum increase at the flare ribbons has a good relationship with ∆B LoS .These confirm that the continuum emission can be linked to ∆B LoS but in an unknown way (Castellanos Durán and Kleint, 2020). 5.The Hα Doppler velocity at the flare ribbons has good correlations with B LoS as well as Hα asymmetry in this small WLF.
SOLA: ms.tex; 3 May 2024; 1:16; p. 10 From a statistical perspective (Section 3.5), we find a good correlation between the CHASE continuum increase and photospheric ∆B LoS at flare ribbons, while the case is not between the WST continuum increase at 3600 Å and ∆B LoS .This may be explained by the different formation heights of these two continua.The CHASE continuum near the Fe i line at 6569.2 Å can be formed lower (say in either the photosphere or chromosphere: e.g.Kleint et al. 2016;Kerr 2017) than the WST 3600 Å continuum (i.e.Balmer continuum), the latter of which mainly originates from the chromosphere during a flare (e.g.Avrett, Machado, and Kurucz, 1986).As regards the Hα line, it is formed from the photosphere to chromosphere from its wing to core.In particular, this line is sensitive to nonthermal heating (e.g.Henoux, Fang, and Gan 1993;Rubio da Costa et al. 2016;Capparelli et al. 2017).Therefore, its Doppler velocity could reflect chromospheric condensation and probably also evaporation (e.g.Song et al., 2023) caused by nonthermal electron beams.The Hα velocity shows some correlation with the photospheric B LoS .This implies that the chromospheric condensation/evaporation can be influenced by the magnetic field in the lower atmosphere particularly in WLFs, as in a WLF, the position of chromospheric condensation/evaporation could be relatively lower in height due to a lower energy deposition height compared with non-WLFs.This small C2.3 WLF has some properties similar to those of large or major WLFs.(1) The CHASE Fe i and Hα lines at the white-light kernels in this small WLF display similar profiles, i.e. symmetric Gaussian profiles still in absorption in Fe i and asymmetric emission profiles with a central reversal in Hα, to those in an X1.0 WLF as reported by Song et al. (2023), although their relative increases in the intensity are much lower compared with the X1.0 WLF.(2) The whitelight emissions in this small WLF are related to a nonthermal electron-beam heating, which is the same as the case in many large WLFs (e.g.Kuhar et al., 2016;Lee et al., 2017;Song et al., 2023).(3) Some white-light brightening points in this small WLF are accompanied by a photospheric ∆B LoS while some are not.This is consistent with the results in X-and M-class flares (Castellanos Durán and Kleint, 2020).All these may hint that the small WLFs with C and even B class may be the same as the large WLFs with M and X class in nature.
More small WLFs are worthy of study in the future, especially with comprehensive observations.Firstly, combining with HMI continuum images at Fe i 6173 Å (in the Paschen continuum) and WST images at 3600 Å (in the Balmer continuum), we can investigate the types of WLFs, i.e. type I or II with or without a Balmer jump (e.g.Fang and Ding, 1995), as well as the origins of these continuum emissions in small WLFs.In particular, WST provides 3600 Å images with a very high cadence of one second in a burst mode, which can capture some rapid variations of white-light brightenings during a WLF.Secondly, using the spectroscopic observations of Fe i and Hα lines from CHASE, we can obtain the dynamics from the photosphere to chromosphere induced by the energy deposition of a WLF.The Doppler velocity and also line width from conjugate footpoints deserve a further study, which could provide some insights into the heating or energy deposition of a WLF.Thirdly, HXR spectral and imaging observations from HXI and/or STIX are necessary to determine the heating mechanisms of a WLF.Some quantitative relationships between SOLA: ms.tex; 3 May 2024; 1:16; p. 11 the HXR flux (or even low-energy cutoff or spectral index of the spectrum) and continuum increase further help to establish a detailed relation between the electron-beam heating and white-light emissions in observations.Finally, magnetic field observations from HMI and also the Full-disk vector Magneto-Graph (FMG: Su et al., 2019a) on ASO-S are important to understand how the magnetic field influences the white-light emissions, which has been unknown even in large WLFs.Overall, on the one hand, a statistical study for small WLFs is worthwhile to be performed using high-resolution observations.On the other hand, with the aid of sophisticated radiative hydrodynamic simulations, some case studies are required to understand the energy transportation and deposition as well as emission mechanisms for both large and small WLFs.
loops are overplotted on AIA 131 Å, WST 3600 Å, and CHASE continuum and Hα images (Figures2a -d).At an earlier time of 04:09:22 UT (Figures SOLA: ms.tex; 3 May 2024; 1:16; p. 7 HMI ∆B LoS .Figures6a -cshow the scatter plots of CHASE continuum increase versus photospheric B LoS , ∆B LoS , and Hα asymmetry.Note that the data points in the shaped area, i.e. within uncertainties, are excluded to further calculate the linear Pearson correlation coefficient (cc).It is seen that the CHASE continuum increase only has some correlation with ∆B LoS (cc=0.52).In contrast, it has no obvious relationship with B LoS (cc=−0.01)orHαasymmetry(cc=0.25).Figures6d -fdemonstrate the scatter plots of WST continuum increase versus B LoS , ∆B LoS , and Hα asymmetry.We find that the continuum increase at WST 3600 Å has no linear relationships with B LoS (cc=−0.28),∆BLoS(−0.23), or Hα asymmetry (cc=0.32).Figures6g -iexhibit the scatter plots of Hα Doppler velocity versus B LoS , ∆B LoS , and Hα asymmetry.We can see that Hα Doppler velocity has good relationships with B LoS (cc=−0.56)as well as Hα asymmetry (cc=0.93).By contrast, it has no good correlation with ∆B LoS (cc=−0.36).Note that the negative correlation between Hα velocity and B LoS is owing to negative magnetic fields at most ribbon locations.The implications or explanations of these relationships among different parameters will be discussed in the following Section.