A MULTI-INSTRUMENT ANALYSIS OF A C4.1 FLARE OCCURRING IN A δ SUNSPOT

The SST (Scharmer et al. 2003) acquired high-resolution data on 2011 August 6 from 09:00:05 UT until 09:37:37 UT, during the gradual phase of the C4.1 flare, when the AR was at solar coordinates (−350'', −360''), with heliocentric cosine angle $\mu =0.84$.

The CRisp Imaging SpectroPolarimeter (CRISP, Scharmer et al. 2008) carried out spectroscopic measurements along the profile of the Fe i line at 557.6 nm and full Stokes profile measurements of the Fe i pair at 630.2 nm. Filtergrams in the core of the Ca ii H line at 396.88 nm, with a passband of 110 pm, and in the adjacent wide band were acquired simultaneously with the Fe i data by the cameras along the blue channel of the SST. The pixel size of the Fe i data was 0059 pixel⁻¹ at 557.6 nm, while that of Ca ii H filtergrams was 0034 pixel⁻¹. The temporal cadence of each complete scan of the Fe i lines was ∼28 s. As concerns the blue beam (Ca ii H line core and wide band), the temporal cadence was 8.9 s. The field of view (FOV) of these SST data sets is $57\_\_AMP\_\_farcs;5\times 57\_\_AMP\_\_farcs;3$.

For further details on the SST data sets, we defer the reader to Cristaldi et al. (2014), who provide more information about the spectral coverage of the SST/CRISP measurements and the data reduction procedure via the CRISPRED pipeline (de la Cruz Rodríguez et al. 2015). Here, we only remind the reader that the application of the multi-object multi-frame blind deconvolution (MOMFBD, van Noort et al. 2005) technique to these SST data sets ensured that near diffraction-limited spatial resolution was achieved (015 at 557.6 nm).

Full-disk continuum and longitudinal (line-of-sight, LOS) magnetograms taken by the Helioseismic and Magnetic Imager (HMI, Scherrer et al. 2012) on board the SDO (Pesnell et al. 2012) in the Fe i line at 617.3 nm with a resolution of 1'' were used to complement the high-resolution data set of the SST. The SDO/HMI data used in this analysis cover three days of observations, starting from 2011 August 5 at 00:10:25 UT until August 7 at 23:58:25 UT, with a cadence of 12 minutes (filtergrams constructed from the HMI vector magnetic field series). All the SDO/HMI images were aligned, taking into account the solar differential rotation, by using the IDL SolarSoft package.

In order to analyze in more detail the magnetic configuration of the AR during the flare, we also benefited from SDO/HMI Space-weather Active Region Patches (SHARPs) data (Hoeksema et al. 2014). These SHARP data provide maps of the photospheric vector magnetic field of the AR and its uncertainty, computed using the Very Fast Inversion of the Stokes Vector code (VFISV; Borrero et al. 2011). The continuum intensity, Doppler velocity, and LOS magnetic field are also provided. See Bobra et al. (2014) for a comprehensive explanation of the SHARP pipeline. We corrected the SHARP data for the rotation angle of 180° of SDO/HMI data, and the vector magnetic field components were transformed into the local solar frame according to Gary & Hagyard (1990) using the proper SDO/HMI procedure (Sun 2013). Finally, we selected a subFOV of these data of about $129^{\prime\prime} \times 91^{\prime\prime}$ encompassing the AR 267.

Data taken by the Atmospheric Imaging Assembly (AIA, Lemen et al. 2012) aboard the SDO mission were used to study in detail the evolution of the flare in the coronal and upper chromospheric layers. We extracted a series of cutout images with a FOV that covers $129^{\prime\prime} \times 91^{\prime\prime}$, the same used for the SHARP data. These SDO/AIA cutouts comprise the time interval between 08:00 UT and 10:00 UT on August 6, with the highest available cadence (12 s for the EUV passbands, 24 s for the UV 1600 Å images).

The XRT (Golub et al. 2007) aboard the Hinode satellite (Kosugi et al. 2007) was observing nearly the full solar disk, including AR 267, through the Ti/poly filter. The Hinode/XRT had been operating in patrol mode since 08:17 UT. During the flare peak at 08:47 UT, it activated the target-of-opportunity mode and acquired a short series of filtergrams at higher resolution in four filters with a cadence of 22.2 s. The Hinode/XRT acquired 4 filtergrams through each of the Ti/poly and Al/mesh filters, with a pixel scale of 206, and 4 filtergrams through each of the Al/thick and Be/thick filters, with a pixel scale of 103. For the latter filters, the exposure times are reported in the Table 2.

The RHESSI (Lin et al. 2002) satellite was in night until 08:59 UT and after this time performed some measurements in the nine different energy channels. One event in the flare list of RHESSI, which has been attributed to AR 267, has two peak emissions occurring at 09:00:06 UT and at 09:11:38 UT. Both of these peaks fall within the SST observing time. In particular, the first peak at 09:00:06 UT, mostly visible in the low energy channels, i.e., at 3–6 and 6–12 keV, was by far more intense than the second one.

SDO/AIA and SDO/HMI data were aligned with each other by using the IDL SolarSoft mapping routines to take into account the different pixel sizes. SST Ca ii H and SST/CRISP images were aligned by considering the first filtergram in the adjacent wide band of the Ca ii H and the continuum image deduced from the first sequence taken by SST/CRISP in the Fe i 557.6 nm line acquired at 09:00:05 UT, which was already aligned with the SST/CRISP data in the Fe i 630.2 nm pair by means of the MOMFBD algorithm.

To co-align SST/CRISP and SDO observations, we used the first spectral image in the sequence of SST/CRISP data taken in the continuum of the Fe i 557.6 nm line and the SDO/HMI reference image in the continuum closest in time. The displacement between the two images was obtained with cross-correlation techniques. The alignment obtained was confirmed by the correspondence between the flare ribbons seen in the SST Ca ii H image taken at 09:00:05 UT and the nearly simultaneous SDO/AIA image acquired at 1600 Å. We estimate that the precision of the alignment obtained by using this procedure is of the order of the pixel scale of SDO images, i.e., 05.

The Hinode/XRT filtergrams were aligned to the SDO/AIA images by using the first non-saturated filtergrams obtained through the Al/thick filter and the simultaneous SDO/AIA image acquired at 94 Å at the flare peak. The displacement between these two images was obtained with cross-correlation techniques. The precision of such an alignment can be estimated as of the order of the resolution of these Hinode/XRT measurements, about 2''. Finally, for RHESSI measurements, we used the pointing information provided by the instrument.

To have diagnostic information about the flare local magnetic environment and the properties of the energy release site, we have analyzed observations taken around the flare peak.

A general survey of the magnetic configuration of the AR 267 in the photosphere during the flare is indeed useful to infer the vector magnetic field topology which could lead to the flare. Figure 3 (top panel) shows the SDO/HMI SHARP map of the longitudinal component of the vector magnetic field nearly at the time of the flare peak. The red (green) arrows indicate the direction of the transversal component in the positive (negative) polarity and their orientation follows the azimuth direction. A strong magnetic shear is detected in the δ spot along the PIL separating the two umbrae of opposite polarities (thick orange line), as the magnetic field is almost parallel to this PIL. The average value of the horizontal shear angle (see, e.g., Gosain & Venkatakrishnan 2010) along the δ-spot PIL is about 80°. The shear angle was calculated considering a potential extrapolation according to Alissandrakis (1981). Diffuse, weak magnetic fields with opposite polarities are observed at the eastern and western edges of AR 267.

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In order to obtain hints about the connectivity of the magnetic field lines of AR 267 in the upper atmospheric layers, we have computed a linear force-free field extrapolation, again using the method of Alissandrakis (1981). We have chosen a value of the α parameter equal to 00025⁻¹, which seems to reproduce in a better way the loops observed in the corona in the EUV (see Section 3.2). Such a force-free field approximation can be considered fairly representative of the overlying field, which does not appear to be strongly twisted (e.g., Régnier 2013). Although in the presence of magnetic field advection and cancellation at the photospheric level, in this manner we can have some information about the global field of AR 267. In Figure 3 (bottom panel) we display an image representing the result of this extrapolation. We find the presence of three different systems of magnetic field lines in AR 267: blue lines represent the field lines connecting the opposite polarities of the δ spot, red lines are those connecting the positive polarity of the δ spot with the following sunspot, and yellow lines are those connecting the positive polarity of the δ spot with the diffuse field of negative polarity to the south of the main negative polarity of AR 267.

The careful analysis of the GOES-15 flux curves in the 1–8 Å channel, shown in Figure 4 (left panel, solid line), indicates that the C4.1 flare peaked at 08:47:37 UT, which is the time corresponding to the change in the sign of the derivative of the soft X-ray (SXR) flux. Moreover, we take the derivative of the SXR flux as proportional to the HXR flux, referring to the so-called "Neupert effect" (Neupert 1968). Looking for an increase in the HXR flux, we have estimated the beginning of the impulsive phase, which occurs at 08:37:51 UT. In addition, in Figure 4 (left panel) we display the GOES-15 flux curves in the 0.5–4 Å channel (dotted line), which shows a similar temporal behavior.

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Figure 4 (left panel, circles) also shows the light curve deduced from the Hinode/XRT measurements through the Ti/poly filter: it has been obtained by taking the average of the X-ray intensity within the box indicated with a black line in the Hinode/XRT filtergrams displayed in Figure 4 (right panel). The quiet Sun regions indicated with gray boxes have been used to verify the genuineness of the X-ray signals that were not due to cosmic-ray contamination. The X-ray flux in these quiet Sun regions is negligible compared to the scale used for Figure 4 (left panel).

When Hinode/XRT observed AR 267 at high resolution during the flare peak, there were filtergrams that did not go into saturation, although they were only those acquired through the Al/thick and Be/thick filters. This fact makes these Hinode/XRT data unique, as very few flares observed by the instrument have exposure times such that only the emission core is visible. All of the four images taken with the Hinode/XRT Be/thick filter were not saturated, while in the Al/thick filter only two images over the four filtergrams did not reach the saturation level, i.e., the ones with shorter exposure times (see Table 2).

Thus, in order to deduce the coronal temperature and the differential emission measure (DEM) we have summed over the two Al/thick images not saturated and the ratio has been calculated for the two Be/thick images closest in time. Indeed, we take advantage of the procedure proposed by Narukage et al. (2011) to estimate the temperature and DEM by means of the filter ratio method using the SolarSoft routine xrt teem.pro. We adopt the calibration performed by Narukage et al. (2014) for those Hinode/XRT thick filters.

In Figure 5 (top panel) we display the X-ray intensity with the Al/thick filter at the flare peak. We can note that the bulk of the X-ray emission is located in the area corresponding to the δ-spot region ([−355'', −335''] × [−365'', −345'']). The plasma temperature is shown in Figure 5 (bottom panel): the average value in the flaring region is about $1.9\cdot {10}^{7}$ K, in agreement with predictions of flare models (see, e.g., Hirayama 1974). The mean value of the volume emission measure in the same region is $46.5\;\mathrm{log}\;{\mathrm{cm}}^{-3}$.

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In this section of the analysis, we study the timing and spatial evolution of the structures as seen by SDO/AIA in the upper atmospheric layers to describe the development of the flare. We have found an asymmetric evolution of the ribbons and the formation of a Y-shaped configuration whose tail expands with time over the following sunspot of AR 267.

In Figure 6 we report a series of snapshots relevant to the C4.1 flare observed by SDO/AIA at different EUV wavelengths around the flare peak at 08:47 UT and almost simultaneous to the SDO/HMI SHARP map shown in Figure 3 (top panel). These images show the morphology of AR 267 at that time in various atmospheric layers. In the upper photosphere we are able to distinguish a pair of asymmetric ribbons (see the image at 1600 Å), while in the upper chromosphere we also detect some emission in between them (see the image at 304 Å). In the corona the emission above the flaring region reaches the saturation level of the detectors for the majority of the EUV channels, i.e., at 171, 193, and 211 Å. Only at 335 and 94 Å  can we note that the ribbons visible in the lower atmospheric layers are overarched by coronal loops that connect them. At these wavelengths, the shape of AR 267 seems very similar to the one simultaneously observed in the X-rays by Hinode/XRT with lower spatial resolution (compare Figure 6, at 335 and 94 Å, with Figure 5, top panel). Note that, as far as this study is concerned, the saturation present in almost all of the EUV channels does not complicate our analysis.

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Interestingly, an extruding structure directed from the flaring region toward the eastern direction is seen at various wavelengths: it is extremely evident in the image in the 131 Å passband, which during flares samples emission from mainly Fe xxi lines (see, e.g., O'Dwyer et al. 2010). The presence of this feature gives rise to a Y-shaped structure in the corona above AR 267, which will become clearer in the subsequent minutes.

We have computed the light curves for all the SDO/AIA EUV channels and for the 1600 Å filtergrams, by averaging the SDO/AIA intensity within the box shown in Figure 6 in the image relevant to the 304 Å channel (top-left panel). These light curves, which are displayed in Figure 7, show that most of the EUV and UV intensities begin to increase at the time of the HXR emission, indicated with a dotted vertical line. The plot also points out the presence of some time delays between the intensity peaks in the different EUV channels between themselves and between the SXR emission measured by the GOES satellite as the SDO/AIA intensities reach their maximum value at slightly different times.

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In detail, the maximum intensity in the lower atmospheric layers, to which the 304 and 1600 Å channels refer, is reached about 30 s before the flare peak in SXR, which is indicated with a solid vertical line in Figure 6. The maximum is almost simultaneous with the flare peak at 131 Å, the EUV filter relevant to the hottest plasma temperature. Moreover, at this wavelength the light curve is narrower than the others and exhibits a linear rise and decrease, the latter with a slope slightly steeper than the former. A similar linear trend is present in the light curve at 94 Å, which, however, has a maximum about 60 s after the flare peak. The light curve at 335 Å has its maximum 120 s after the flare peak with a slow decrease. Finally, the 171, 193, and 211 Å channels, which are strongly affected by saturation, exhibit a rapid rise phase, with a maximum almost simultaneous with the flare peak, and have a very long decrease phase. Table 3 summarizes the time delays that are found between the intensity peaks in the light curves relevant to different EUV passbands, between them and the SXR emission.

To study the spatial evolution of the flare ribbons in the δ-spot region, we have used intensity time slices. In the map shown in Figure 6 relevant to the 304 Å channel (top-left panel), we have indicated three segments with different colors labeled with different letters (A, B, C), along which the intensity time slices have been derived. We show in Figure 8 the slices relevant to four SDO/AIA passbands that exhibit a low degree of saturation: 335, 211, 1600, and 304 Å. In these images, the bottom part of the slice refers to the southwestern point of the segments in Figure 6 and the top part to the northeastern point of the segments.

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The activation of the emitting regions, which form the ribbons, occurs first in segment A even if an emission kernel is present since the beginning of HXR emission in the northern ribbon along segment C. Then the emission in ribbons activates in segment B and eventually in segment C. The northern ribbon activates before the southern one. The ribbons' activation also shows a delay between the atmospheric layers: it begins in the 1600 Å channel, then it is seen at 304 Å, at 211 Å, and finally at 335 Å. However, the southern ribbon along segment C activates simultaneously in all of the SDO/AIA channels, almost at the time of the flare peak at 08:47 UT.

The flare ribbon motions are clearly recognized at 1600 and 304 Å, along the segments B and C. We note both the effect of the drift of the northern ribbon outward of the δ-spot region, and the similar motion of the southern ribbon, when it appears, which results in the increasing flare ribbon separation. Note that the different slope in the drift of the ribbons indicates that the southern ribbon, when it appears, moves faster than the northern ribbon. The estimate of the speed, which has been derived from the 1600 Å observations through a linear fit of the displacements of the intensity maxima in the ribbons, is of 1.05 and $1.75\;\mathrm{km}\;{{\rm{s}}}^{-1}$ for the northern and southern ribbons, respectively.

Conversely, the emission in ribbons has already ended along segment A at 1600 Å at the time of the flare peak. After that time, in the upper chromosphere the bulk of emission ends along segment A, then in segment B, and finally in segment C (see the 304 Å passband). A similar progression is present in the upper atmospheric layers, as can be seen in the images relevant to the 211 and 335 Å channels, but at these levels the emission decrease is slower than in the chromosphere.

These dynamics are highly suggestive of an asymmetric ribbon expansion. Besides separating each other, as usual, the flare ribbons move toward the northeastern part of the δ-spot region along a direction that is perpendicular to that identified by the crossing segments A, B, and C shown in Figure 6.

We have also derived the evolution of the extruding structure in all the UV/EUV channels using a slice along it as shown in Figure 9. This figure shows the map of AR 267 in the 131 Å passband at the flare peak. The red line indicates the slice that has been used for the subsequent analysis.

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The behavior of the extruding structure in the various wavelengths is reported in Figure 10. In this figure, the bottom part of each panel refers to the easternmost point along the slice, while the top part of the panels refer to the point of the slice closest to the δ spot. We can see that this structure first appears at 131 Å, where it is observed for less than ten minutes since 08:44 UT, a few minutes before the flare peak (indicated in this panel with a black vertical line). Slightly after the flare peak, when the EUV intensity reaches its maximum at 131 Å, the structure is also observed at 94 Å: in this channel the emission is no longer detected after 09:00 UT. In both these channels, the emission kernel is at first observed near the δ spot. Then, the EUV intensity peak moves toward the following spot of AR 267, at wavelengths progressively referring to lower plasma temperature. The structure is observed at 335 Å about five minutes after the flare peak, and the emission enhancement lasts for more than 30 minutes. We also note a very faint enhancement near the δ spot at 08:47 UT, at the flare peak, and in the 193 and 171 Å channels, which can be also found at 304 and at 1600 Å. Much more evident is the sudden increase in brightness that is simultaneously detected at 211, 193, and 171 Å at 08:57 UT, corresponding to the appearance of the extruding structure in these channels. This EUV emission increase is observed mostly near the following spot of the AR and it is dramatically strong at 171 Å for a couple of minutes. After about one minute, an abrupt dimming at the chromospheric level is detected at 08:58 UT along the slice in the 304 Å passband, which lasts for a few minutes. Other enhancements are observed at 09:06 UT and 09:10 UT at 193 and 171 Å, which are located mostly near the following spot. Near the δ spot a new enhancement along the slice is seen at 09:10 UT in the 304 Å passband with a short duration. The extruding structure is not found at 1600 Å where, except for the brightening at the time of the flare peak, when we notice only slight enhancements toward the δ spot and the following spot of AR 267.

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The evolution of the extruding structure has been further investigated by considering the average value of the EUV intensity in four passbands in a region encompassing the structure, framed by a solid box in Figure 9. The resulting light curves are displayed in Figure 11. The EUV intensity has been normalized to the maximum value of each passband in the time interval considered in the plots.

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The intensity in the 131 Å passband has a maximum at the time of the flare peak, and only a slight increase during the decay phase, almost coincident with the maximum seen at 211 Å, at 08:58 UT. A few minutes before, we find the maximum in the 335 Å channel, after a slow increase. The increase and the immediately subsequent dimming are clearly visible in the 304 Å at 09:00 UT, as indicated by the arrow in Figure 11. Such a dimming is followed by a new, slight increase and finally the intensity decreases.

In order to discuss the spatial evolution of the flare in different SDO/AIA passbands during the early gradual phase, we display in Figure 12 the isophotes of the EUV/UV emission at different times. Contours drawn in red color at 08:46 UT, at the flare peak, show strong, localized emission straggling over the δ-spot PIL (solid orange line). This emission reaches saturation levels in some EUV passbands. The flaring site is still localized, even if with a lesser extent over the δ-spot PIL at 08:53 UT (contours shown in green color). This is seen in the various channels, for instance at 211 Å, but is more pronounced in the 304 Å channel where the EUV emission is placed almost at the opposite sides of the δ-spot PIL. At the same time, we see the appearance of the extruding structure in the EUV channels referring to lower plasma temperature. It can be recognized as new emission regions appearing to the south of the main PIL of AR 267 (solid light blue line in Figure 12) almost parallel to it. This is particularly evident in the 211 and 335 Å channels.

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By the time SST observations began, the blue contours in Figure 12 referring to 09:00 UT indicate that the extruding structure was clearly formed and connected the initial flaring region, localized around the δ-spot PIL, to the following polarity of the AR 267. The structure appears to run parallel to the main PIL of the AR: this is distinctly seen at 211 Å. At 335 Å we find that the structure still occupies the same region that was occupied at 08:53 UT, while in the 1600 Å passband we only see a small patch appearing in that region. In the δ-spot region, the EUV emission is now reduced in size, placed almost at the opposite sides of the δ-spot PIL also in the EUV channels referring to higher plasma temperature (335 and 211 Å). In the passband referring to the lower layers of the atmosphere, the EUV/UV emission is localized in small emission patches at the opposite sides of the δ-spot PIL (see the contours at 1600 and 304 Å). Note that the apparent motion of the extruding structure eastward, i.e., of the EUV/UV enhancements, occurs along a direction that is nearly opposite to the direction of the flare ribbon expansion in the region of the δ-spot PIL.

In this section we benefit from high-resolution ground-based observations and RHESSI measurements, which cover only the late gradual phase of the C4.1 flare, since about 13 minutes from the beginning of the flare. We have studied the evolution of the ribbons, their fine structure magnetic environment, and the correlation of their spatial distribution with the presence of HXR sources.

Figure 13 shows the chromospheric appearance of AR 267 in the core of the Ca ii H line at the beginning of the SST observations (reverse intensity scale). It is possible to clearly distinguish three regions with enhanced emission, which correspond to the three apparent ribbons of the flare, encircled with lines.

We have computed the light curves relevant to the three sites with enhanced emission observed in the chromosphere, highlighted in Figure 13. The plots shown in Figure 14 report the total UV intensity (left panel) and the average UV intensity (right panel), normalized for the background represented by the average value of the Ca ii H line core intensity within the rectangular box overplotted in Figure 13. For each of the three regions (see the line code in the caption), we have considered the total and the average value of the UV intensity only for those pixels with emission at least twice the value of the background intensity.

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The two bigger regions, corresponding to the northern ribbon in the δ-spot region and to the ribbon in the following sunspot, respectively, exhibit a decreasing trend during the gradual phase for both the total intensity and the average intensity, as usually observed during the gradual phase. In contrast, we note that the small region corresponding to the southern ribbon in the δ spot shows an almost constant value in the total intensity, while the average intensity has some enhancements during the decrease. The strongest enhancement seems to be simultaneous with the second peak of X-ray emission observed by RHESSI at 09:11:38 UT. A similar spiky behavior, characterized by a timescale of a couple of minutes, can also be seen in the other two ribbons, concerning more clearly the total UV intensity. However, this can be in some measure due to variations in the seeing or to oscillations in the chromosphere (e.g., Jess et al. 2015) as it partially affects the background emission measured in the rectangular box of Figure 14.

SST/CRISP observations provide new information concerning the photospheric configuration of AR 267 during the gradual phase. In Figure 15 (top-left panel) the morphology of AR 267 at the beginning of SST/CRISP observations in the continuum of the Fe i line at 603.15 nm is shown. We remind the reader that the characteristics in the region around the δ-spot PIL in the photosphere have been extensively described by Cristaldi et al. (2014). Here, we note that the preceding and the following sunspots appear comparable in size and that light bridges are present in the umbral regions of both sunspots.

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Figure 15 (top-right panel) displays the simultaneous map of Doppler velocity, which shows the plasma motions at a photospheric level in the AR 267. To derive the Doppler velocity, we have applied a Gaussian fit to the Fe i line profile at 603.15 nm with the MPFIT routine (Markwardt 2009) in IDL. The local frame of rest was calibrated by imposing that plasma in the quiet Sun has, on average, the convective blueshift for the Fe i 630.15 nm line (Dravins et al. 1981), corrected for the position on the solar disk ($\mu =0.84$) according to Balthasar (1988). We have used a value of $-235\;{\rm{m}}\;{{\rm{s}}}^{-1}$.

The map shows strong upflows and downflows in the region between the opposite polarities of the δ spot, as observed by Cristaldi et al. (2014) in other iron lines. At the easternmost extremity of the northern ribbon in the δ-spot region, we find an upflow of about $1.5\;\mathrm{km}\;{{\rm{s}}}^{-1}$.

Figure 15 also shows the simultaneous maps of the magnetic field intensity (bottom-left panel) and of the longitudinal component of the magnetic field (bottom-right panel) of the entire SST FOV. These magnetic field maps have been deduced by applying the VFISV code, which performs a Milne–Eddington inversion of the data (Borrero et al. 2011), to the full SST/CRISP spectropolarimetric profiles along the Fe i 630.15 nm line. The longitudinal component refers to the observer frame, and although it does not take into account the effects of projection, it is sufficient for the scopes of our analysis.

The fine structure of the tangled magnetic field in the δ-spot region is clearly revealed in the longitudinal field map. The magnetic field strength is of the same order of magnitude in the preceding and in the following sunspot—about 1500 G. The latter is starting a decay phase, as confirmed by later observations: this is suggested by the fact that a light bridge with a weaker, more inclined magnetic field of about 700 G is splitting the sunspot into two parts.

At the time of the beginning of the SST observations, the RHESSI satellite also performed some measurements over AR 267. The information in the RHESSI map refers to a flare peaking at 09:00:06 UT. The reconstructed images have been made using the CLEAN algorithm (Hurford et al. 2002). In Figure 16 we draw the contours of the X-ray emission measured in three different channels: 3–6 keV (red color), 6–12 keV (green color), and 25–50 keV (black color) over the simultaneous SST/CRISP continuum map. The contours refer to 80%, 90%, and 98% of the maximum of the emission measured by RHESSI for each channel, respectively. The X-ray emission takes place outside of both the sunspots, in the region between them. Thus, at this time it does not seem to be spatially correlated with the δ-spot region.

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We have also produced three maps from the simultaneous SDO/AIA observations at 09:00 UT, overplotting the contours of the X-ray emission measured by RHESSI with the same color scheme as before. Figure 17 displays the SDO/AIA 335 Å map (top panel), the SDO/AIA 1600 Å map (middle panel), and the SDO/AIA 304 Å map (bottom panel). The X-ray emission is localized around the crossing point of the Y-shaped structure (see the 335 Å map) formed by the apparent three ribbons of the flare at that time. The 304 Å map also reveals that the contours of the brightenings that form the Ca ii H flare ribbons observed by SST coincide with the footpoints of the post-flare coronal loop seen in this passband (compare also with Figures 13 and 16). The emission kernels of the flare ribbons are clearly distinguishable in the 1600 Å map.

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The inner contours of the 3–6 and 6–12 keV emission measured by RHESSI are located in the region characterized by the extruding structure previously analyzed, suggesting that HXR emission regions could be moving along it. In this regard, it is worth noting that the second peak observed by RHESSI at 09:11:38 UT, detectable only in the 3–6 keV channel, is localized even closer to the crossing point of the Y-shaped structure.

However, one has to be cautious with the analysis of such RHESSI data. Since the satellite was in night during the C4.1 flare peak, the flare detected at 09:00:06 UT by the RHESSI automatic system is due to the sudden increase of the X-ray signal at the dawn of the satellite, which indeed caught the tail of the X-ray emission after the C4.1 flare. This can be easily seen in the corrected RHESSI observing summary rates referring to this event. Thus, while the site of the RHESSI signal at 09:00 UT is really coincident with the crossing point of the Y-shaped structure seen in the EUV maps, this must be interpreted only as the site of the X-ray emission at the time of the RHESSI observations and does not correspond to the true kernel of the C4.1 flare at the time of its ignition.

The event seems to occur during a magnetic flux decrease phase. This suggests that the flare originated in a region where magnetic cancellation was taking place, likely due to the rearrangement of the magnetic field in the δ-spot region (see, in particular, the red and blue symbols in Figure 2, middle panel, slightly after the flux peak value, $\sim 1.5\times {10}^{21}$ Mx, just before the C4.1 flare).

We notice from Figure 2 (bottom panel) that the magnetic helicity accumulation is increasing—in absolute terms—during the evolution of the AR 267, reaching a value of $\sim -7.3\times {10}^{40}\;{\mathrm{Mx}}^{2}$ at the time of the C4.1 flare from the beginning of the observations. Moreover, the helicity of the AR 267 is in contrast with the general cycle-invariant hemispheric helicity rule.

The violation of the hemispheric helicity rule (Liu et al. 2014), the tangled fine structure of the longitudinal field in the δ-spot region and, above all, the presence of a large shear angle ($\sim 80^\circ$) along the δ-spot PIL are clear indications of the magnetic complexity of the AR 267. It is worth noting that Cristaldi et al. (2014) found strong, persistent motions of both upflows and downflows along the δ-spot PIL, which seemed not to be related with the occurrence of the C4.1 flare, and evidence of sheared magnetic field lines in the same region.

In this regard, Shimizu et al. (2014) also have recently reported on a long lasting high-speed material flow observed in a δ-spot PIL in AR NOAA 11429. They found that this plasma flow began six hours before the onset of a X5.4 flare, along the PIL located between the flare ribbons, and continued for several hours after the flare. They interpreted these observations as material flow increasing the magnetic shear along the δ-spot PIL, also being able to develop a magnetic structure favorable for the triggering of the eruptive flare.

Hence, we cannot exclude that the motions observed by Cristaldi et al. (2014) in the Fe i 557.6 and 630.25 nm lines, which we also found in the 630.15 nm line in the same area between the opposite polarities of the δ spot, accumulate magnetic shear. When the latter is enough to trigger restructuring of the field lines via reconnection, the C4.1 flare occurs.

The magnetic field configuration inferred from the linear force-free field extrapolation (see Figure 3, bottom panel) indicates that at the coronal level there are different bundles of magnetic field lines, characterized by different connectivities: those connecting the two magnetic polarities inside the δ spot (blue lines in Figure 3, bottom panel), those connecting the positive polarity of the δ spot with the following sunspot (red lines in Figure 3, bottom panel), and those connecting the positive umbra of the δ spot with the southernmost, most diffuse negative polarities (yellow lines in Figure 3, bottom panel). Note that, even if this linear force-free field extrapolation is only an approximation, the topology of the AR field is qualitatively robust. This configuration suggests that any rearrangement involving the blue field lines will perturb the other systems of field lines, as will be discussed to a deeper extent in Section 5.

The SXR/HXR emission measured by the GOES-15 satellite indicates the beginning of the impulsive phase of the flare at 08:37:51 UT, with the flare peak at 08:47:37 UT. The Hinode/XRT emission is in agreement with this flare timing.

Time delays are found between the intensity peaks in the light curves relevant to different EUV passbands: in particular the SDO/AIA channels relevant to the lower atmospheric layers (304 and 1600 Å) have a maximum before the SXR flare peak, which in turn is simultaneous with the peak in the hardest 131 Å filter. On the other hand, the analysis of the SDO/AIA time slices relevant to the region of the two ribbons above the δ-spot indicates that the first signatures, i.e., brightenings, are observed in the corona, at 211 Å, and, although less evident, also in the transition region and chromosphere, at 304 Å. These signatures could be indicative of a not very high atmospheric location of the initial site of energy release, maybe not far from the chromosphere (Falchi et al. 1997).

Concerning the time slices relevant to the so-called extruding structure, Figure 10 indicates that initially this site appears to brighten at 1600 Å in the region close to the δ spot, and later on it becomes brighter along its entire length at coronal heights: at 94 Å slightly before the SXR peak, and 5 minutes after this it is observed at 335 Å. However, subsequent episodes of enhancement and dimming are observed in the corona, i.e., in the 211, 193, and 171 Å passbands in the region close to the following sunspot, and also at a lower height in the region close to the δ spot, as visible at 304 Å (see also Figure 11). From the inspection of Figure 10, these episodes have a timescale of the order of 1–2 minutes. A similar behavior of rising and decreasing of the intensity on the same timescale is observed in the Ca ii H light curves (see Figure 14). Although fluctuations in the seeing or oscillations in the chromosphere can be partially responsible for this effect, the presence of transient enhancements suggests that a mechanism able to cause intensity variations also may act at the chromospheric level. However, we have no hints whether this mechanism is due to a pure chromospheric phenomenon. In any case, this finding could provide a useful constraint for flare models. Indeed, radiative-hydrodynamic codes like RADYN (Carlsson & Stein 1997) should mimic the behavior of these Ca ii H light curves once calibrated in absolute values by means of the comparison with MHD simulations for different combinations of energy input, power-law indices, and energy cut-offs.

Hinode/XRT indicates that at the flare peak the bulk of X-ray emission is located in the area of the δ spot. Based on the Hinode/XRT measurements, the estimate of the plasma temperature at the flare peak is $\approx 1.9\cdot {10}^{7}$ K; the volume emission measure, $46.5\;\mathrm{log}\;{\mathrm{cm}}^{-3}$.

The EUV emission recorded by SDO/AIA shows a different morphology of the AR 267 at the flare peak in the different wavelengths, as well as a different evolution during the gradual phase.

The intensity time slices indicate an asymmetric behavior between the northern and the southern ribbons observed in the δ spot regions: the northern ribbon activates before the southern one, which conversely moves faster when it appears. Moreover, the ribbon activation shows a time delay between the atmospheric layers. Interestingly, the flare ribbons' motions indicate a preferential direction: the separation of the flare ribbons progresses along the northeastern direction of the δ spot. These motions imply that the reconnection X-point is likely traveling laterally as well as vertically as the flare erupts.

The spatial evolution in the EUV/UV passbands indicates a strong and localized emission straggling over the δ-spot PIL at the flare peak, encompassing a lesser extent at 08:53 UT, when a structure departing from the initial flaring region is also observed. At 09:00 UT the extruding structure connects the initial flaring region to the following polarity of the AR 267, being parallel to the main PIL of  AR 267. The apparent motion to the east of the intensity peak along the extruding structure occurs along a direction opposite the other flare ribbons, observed in the δ-spot region. The simultaneous presence of the ribbons and of the expanding extruding structure gives rise to a Y-shaped configuration in the corona.

During the gradual phase, the evolution of the EUV/UV emission, obtained from both SDO/AIA and SST Ca ii H observations, indicates that at 09:00 UT, AR 267 shows the presence of three flare ribbons. Two are the "regular" pair of ribbons at opposite sides of the δ-spot PIL (orange in Figure 12) and the other corresponds to the remote brightenings that are found in proximity of the following sunspot. This site coincides with the eastern footpoint with enhanced emission of the extruding structure at the opposite side of the main PIL of the AR (blue in Figure 12) with respect to the pair of δ-spot ribbons.

RHESSI measurements, compared with the SST/CRISP image in the continuum, indicate that at 09:00 UT the X-ray emission mainly takes place outside of both the sunspots, in the region between them. From the comparison of RHESSI measurements with the EUV images in the 335, 1600, and 304 Å, we can infer that at 09:00 UT the X-ray emission is located around the crossing point of the Y-shaped structure.

In our observations, we have noticed that the progression of the activation of the flare ribbons, which separate with time, follows a preferential direction (see Figure 8) parallel to them. Moreover, we have found that the emitting region apparently moves eastward toward the following sunspot of the AR 267 along a direction that is nearly opposite to the direction of the flare ribbon expansion in the region of the δ-spot PIL.

These dynamics are highly reminiscent of the slipping reconnection (Aulanier et al. 2006) observed by Masson et al. (2009). Analyzing observations of a confined flare presenting a circular ribbon, Masson et al. (2009) noted signatures of field lines reconnecting by following a sequential order along a preferential direction. The authors deduced the magnetic configuration of the AR during the flare from a potential extrapolation and found that (i) the reconnection initially takes place for field lines closer to a coronal null point, and that (ii) it eventually progresses further toward the direction of this null point. Thus, Masson et al. (2009) found that the evolution of the reconnection is induced by the asymmetric magnetic configuration.

Similarly, Wang & Liu (2012) invoked such a model of three-dimensional magnetic reconnection in a fan-spine topology to explain their observations of circular flare ribbons as well as Liu et al. (2013), who found non-standard motions of flare ribbons and a dome-shaped magnetic structure of the flaring region.