# Resolving Two Distinct Thermal X-ray Components in A compound Solar Flare

Zhenjun Zhou(周振军) Planetary Environmental and Astrobiological Research
Laboratory (PEARL), School of Atmospheric Sciences, Sun Yat-sen University,
Zhuhai, China  Key Laboratory of Lunar and Deep Space Exploration, National
Astronomical Observatories, Chinese Academy of Sciences, Beijing 100101, China
CAS Center for Excellence in Comparative Planetology, China  Rui Liu CAS Key
Laboratory of Geospace Environment, Department of Geophysics and Planetary
Sciences, University of Science and Technology of China, Hefei, Anhui 230026,
China  Jianqing Sun School of Astronomy and Space Science, Nanjing University,
Nanjing 210023, China  Jie Zhang Department of Physics and Astronomy, George
Mason University, Fairfax, VA 22030, USA  Mingde Ding School of Astronomy and
Space Science, Nanjing University, Nanjing 210023, China  Xin Cheng School of
Astronomy and Space Science, Nanjing University, Nanjing 210023, China  Yuming
Wang CAS Key Laboratory of Geospace Environment, Department of Geophysics and
Planetary Sciences, University of Science and Technology of China, Hefei,
Anhui 230026, China  Xiaoyu Yu  Planetary Environmental and Astrobiological
Research Laboratory (PEARL), School of Atmospheric Sciences, Sun Yat-sen
University, Zhuhai, China  Lijuan Liu Planetary Environmental and
Astrobiological Research Laboratory (PEARL), School of Atmospheric Sciences,
Sun Yat-sen University, Zhuhai, China  Jun Cui Planetary Environmental and
Astrobiological Research Laboratory (PEARL), School of Atmospheric Sciences,
Sun Yat-sen University, Zhuhai, China

###### Abstract

X-ray emission provides the most direct diagnostics of the energy-release
process in solar flares. Occasionally, a superhot X-ray source is found to be
above hot flare loops of $\sim\,$10 MK temperature. While the origin of the
superhot plasma is still elusive, it has conjured up an intriguing image of
in-situ plasma heating near the reconnection site high above the flare loops,
in contrast to the conventional picture of chromospheric evaporation. Here we
investigate an extremely long-duration solar flare, in which EUV images show
two distinct flare loop systems that appear successively along an
$\Gamma$-shaped polarity inversion line (PIL). When both flare loop systems
are present, the HXR spectrum is found to be well fitted by combining a hot
component ($T_{e}\sim$12 MK) and a superhot component ($T_{e}\sim$30 MK).
Associated with a fast CME, the superhot X-ray source is located at top of the
flare arcade that appears earlier, straddling and extending along the long
‘arm’ of the $\Gamma$-shaped PIL. Associated with a slow CME, the hot X-ray
source is located at the top of the flare arcade that appears later and sits
astride the short ‘arm’ of the $\Gamma$-shaped PIL. Aided by observations from
a different viewing angle, we are able to verify that the superhot X-ray
source is above the hot one in projection, but the two sources belong to
different flare loop systems. Thus, this case study provides a stereoscopic
observation explaining the co-existence of superhot and hot X-ray emitting
plasmas in solar flares.

Sun: flares — Sun: corona — Sun: X-rays

††journal: ApJ††journal: APJ

_Online-only material_ : animations, color figures

##  1 Introduction

The creation of hot (10–20 MK) plasma in the corona is a prominent feature in
nearly all solar flares. Plasma at such high temperatures emits soft X-rays
(SXRs) from both bremsstrahlung and resonant lines of highly ionized atoms -
primarily from Fe XXIV and XXV (Korneev et al., 1979; Caspi & Lin, 2010).
During the flare impulsive phase, the light curve of SXR flux often shows a
tendency to resemble that of the time integral of hard X-ray (HXR) flux. This
empirical relationship, which is also known as the Neupert effect (Dennis &
Zarro, 1993), has provided evidence for the thick-target model (Brown, 1971) ,
in which the HXR emission is produced by the bremsstrahlung of energetic
electrons as they are instantly thermalized in the dense chromosphere,
presumably at the footpoints of newly reconnected field lines, which heats up
the local chromospheric plasma to temperatures in excess of 10 MK; the
overpressure of the over-heated chromosphere propels hot plasma upward into
the corona along the same field lines, forming X-ray emitting flare loops
(Antonucci et al., 1984; Fisher et al., 1985; Allred et al., 2005, 2015). This
upward flow is conventionally termed as chromospheric evaporation (Antiochos &
Sturrock, 1978; Cheng et al., 2019), which is often thought to be the
ubiquitous source of the $\sim\,$10–20 MK plasma observed in nearly all
flares, whose temperatures are consistent with those derived from SXRs of the
_Geostationary Operational Environmental Satellite_ (_GOES_). Termed loop-top
source in the literature, an HXR thermal component with temperatures of 10–20
MK is often detected at the top of the SXR flare loops with indirect imaging
methods (Liu et al., 2013; Sun et al., 2014, 2016).

In some flares, besides the ubiquitous loop top hot component (10–20 MK), a
spatially distinct superhot ($T_{e}\geqslant$ 30 MK) thermal component has
been reported. This superhot plasma is first unveiled with high-resolution HXR
spectroscopy (Lin et al., 1981), which is characterized by a steeply-falling
spectrum resembling that of $\sim\,$34 MK plasma. Continuum and Fe XXVI line
observations showed such a superhot component generally exists in GOES X-class
flares (Tanaka, 1987; Pike et al., 1996). Due to the limited spatially
resolved observations of high-temperature passbands, the source region of this
superhot component ($T_{e}\geqslant$ 30 MK) remains elusive. Case studies
offer a glimpse of location of the superhot source region. With the aid of
direct SXR images by Yohkoh (Ogawara et al., 1991), Nitta & Yaji (1997)
reported a flare with the superhot component consisting of two separate loop
structures, with the dominant HXR flux from an extended structure away from
the bright SXR loop. With the aid of HXR imaging and spectroscopy implemented
by _Reuven Ramaty High Energy Solar Spectroscopic Imager(_RHESSI_ ; Smith et
al., 2002; Lin et al., 2002), Caspi & Lin (2010))_, found that, in a GOES X4.8
flare on 2002 July 23, the superhot plasma is located distinctly above the
flare loop top containing the conventional 10–20 MK plasma, peaks
simultaneously as the non-thermal HXRs, and exists even during the pre-
impulsive phase with negligible footpoints. These observations suggest that
the super-hot plasma is in-situ heated, i.e., more directly related to the
accelerated non-thermal electrons and hence to the reconnection process than
the cooler flare plasma, which is due to the traditional picture of
chromospheric evaporation (Caspi & Lin, 2010). Employing HXR spectroscopy to
derive the temperature of overall flare plasma, Caspi et al. (2014) found a
strong correlation between the maximal temperature and the flare GOES class in
37 M-class-and-above flares. But “super hot” temperatures exceeding 30 MK are
found almost exclusively in X-class flares. It is unclear, however, whether
these flares contain both super-hot and hot components like the prototypical
super-hot flares reported before (e.g., Lin et al., 1981; Nitta & Yaji, 1997;
Caspi & Lin, 2010). Nevertheless, our knowledge about the superhot component,
including its spatio-temporal relationship to the energy release and transport
processes that are active within most flares, is still scarce.

Moreover, some flares have a SXR light curve containing two or more peaks,
which are as close as minutes apart. These multi-peaks are often associated
with two or more closely connected magnetic structures erupting consecutively
within a short time interval. Such a flare is also termed a compound eruption
(Woodgate et al., 1984; Dhakal et al., 2018).

In this paper, we investigate a compound flare that lasted for an extremely
long duration and proceeded sequentially in space along a curved PIL. The long
duration and the optimal projection provide us an excellent opportunity to
analyze the thermodynamic evolution in both time and space of the flare. The
flare is well observed close to the limb by _Atmospheric Imaging Assembly_
(_AIA_ , Lemen et al. 2012) filtergram and _Helioseismic and Magnetic Imager_
(_HMI_ , Schou et al. 2012) onboard _Solar Dynamics Observatory_ (_SDO_ ,
Pesnell et al. 2012), _RHESSI_ , and _GOES_. Meanwhile, this flare is also
well observed from an vantage point by Extreme Ultraviolet Imager (_EUVI_ ,
Howard et al. 2008) telescope onboard Solar TErrestrial RElations Observatory
A (_STEREO-A_ , Kaiser et al. 2008), which observed the flare as an “on-disk”
event, with a spatial resolution of 1.6$\arcsec$ and a cadence of 10 minutes.

These multi-wavelength and dual-perspective observations enable us to conduct
a comprehensive study of the involved thermal structures and their dynamic
evolution in this flare. The paper is organized as follows. We present the
observations in Sect. 2. The data analysis and results are described in Sect.
3, followed by a discussion and conclusions in Sect. 4.

##  2 Observations

On 2012 July 17, a GOES class M1.7 flare occurred near the southwestern limb
of the Sun. This flare began at $\sim$12:24 UT and took $\sim$5 hours long to
reach the SXR peak at $\sim$17:15 UT and then took over 7 hours to reach the
pre-flare level, making the whole duration longer than 12 hours (Figure 1(a)).
In contrast, a typical flare lasts from a few minutes to tens of minutes, and
a long-duration flare lasts for hours, also known as long-duration-event (LDE)
flare (Sheeley et al., 1983; Webb & Hundhausen, 1987). It was termed ‘the
slowest flare’ by Sam Freeland and Hugh
Hudson111http://sprg.ssl.berkeley.edu/$\sim$tohban/wiki/index.php/The Slowest
Flare. In addition, we also measure the time derivative of GOES SXR flux df/dt
during the solar flare (see red line in Figure 1(a)), which can be used as a
proxy for the HXR flux according to the Neupert effect (Dennis & Zarro, 1993).
It should be noted that RHESSI HXR emission (usually defined as X-ray emission
above $\sim$20 keV) is not favored for this extended flare because of the
frequent gaps and low count rates. According to the three peaks in the time
derivative of GOES SXR (red profile in Figure 1(a)), the rising phase of the
flare can be further divided into three episodes: Episode I, from $\sim$12:24
UT to $\sim$14:05 UT; Episode II, from $\sim$14:05 UT to $\sim$16:10 UT; and
Episode III, from $\sim$16:10 UT to $\sim$17:35 UT. The first and last
episodes were associated with two coronal mass ejections (CMEs) at $\sim$
13:48 UT and $\sim$17:00 UT, respectively. Both CMEs launched towards the
southwest (manifested by the central position angle (CPA), which is measured
counterclockwise from the projection of the Sun’s north pole of the broadside
CMEs). The first was a fast CME with a velocity of 958 km/s while the latter
was a slow CME with a velocity of 395 km/s (See the CME height-time plots in
Figure 1(b)), implying the whole process was composed of at least two
different loop systems instead of a superposition of an extended sequence of
similar loops along the PIL.

This extremely long-duration flare also spanned a large area in space,
approximately 250$\arcsec$ along the south-north direction in the NOAA active
region (AR) 11520. The AR is characterized by a major sunspot of positive
polarity surrounded by diffuse magnetic flux of negative polarities. As a
result, the flaring PIL takes a $\Gamma$ shape, with the long ‘arm’ in the N-S
orientation and the short arm in the E-W orientation (Figure 2(e–g)). To
investigate the spatial locations of the thermal components, we reconstruct
the RHESSI X-ray sources in the energy bands of 6-25 keV, the integration time
of the images is 40s. We use the standard image reconstruction CLEAN
algorithm. The CLEAN method is an iterative algorithm (Hurford et al., 2002).
It is widely employed in X-ray image reconstruction due to the excellent
record of bringing out the X-ray image morphology. It is basically a process
of “deconvolution” of the back-projected image using the point spread function
(PSFs). Detectors 3-9 are used, but without detector 4 because it is
excessively noisy during this flare. For all other parameters, the defaults
are used.

In different episodes, the flare showed distinct emission structures at
different locations. In the first episode, there appeared a group of sheared
loops in the AIA 131Å passband in the southernmost part of the active region
(Figure 2(a) and accompanying animation). These loops were relatively low in
height with one compact footpoint patch rooting in the positive magnetic
polarity (northwestern part) and one extended footpoint patch rooting in the
negative magnetic polarity (southeastern part), where the corresponding
brightenings were clearly observed in AIA 1600Å passband (Figure 2(e)). Above
the loops seen in EUV, there existed a 6-25 keV X-ray loop-top source,
implying that these EUV loops were likely hot post-flare loops produced by
magnetic reconnection in the corona. During the one and half hours long
evolution of the first episode, the morphology and the height of the EUV loops
did not show significant changes. However, the SXR emission intensity kept
increasing. Besides, the flare ribbon in the negative polarity showed a
considerable separation movement from the PIL and a northward expansion
(Figure 2(f)). Above the stationary and low-lying EUV loops, there existed a
faint large-scale loop-like structure visible in 131 Å passband from 13:10 UT,
which slowly rose for about 30 minutes and quickly erupted after 13:43 UT.
This eruption resulted in a large CME seen in coronagraph images (black plus
symbols in Figure 1(b)).

Following this quick eruption, the flare evolved into the second episode and
showed a sequential evolution in space along the PIL from south to north
(Figure 2(b) and (f)). The loop-top X-ray source in RHESSI became highly
extended in the N-S direction. In addition, the flare ribbon in the negative
polarity region quickly expanded northward in association with the sequential
formation of the post-flare loop arcade seen in the 131 Å passband. The flare
ribbon initially developed in parallel with the PIL, then showed certain
separation perpendicular to the PIL (Figure 2(f)), probably due to the ascent
of the magnetic reconnection site. The second CME launched at around 16:20 UT
(red plus symbols in Figure 1(b)). After that, a new group of post-flare loops
became visible at $\sim$16:28 UT (Figure 2(c)), transiting into the third
episode of the rising phase.

During the third episode, the flare arcade develops along the E-W oriented PIL
segment toward the limb (Figure 2(d) and (g)). Thus, it became more difficult
to observe the flare evolution along the PIL due to projection effects.
Incorporating the STEREO observation (Figure 5(b) and accompanying animation),
one can see that the increased brightness mostly came from a compact region at
the northernmost part.

##  3 Analyses and Results

The analysis of thermal properties of flare regions can help us infer where
the energy is released from magnetic reconnection. Through tracking the
evolution of thermal sources, we can also infer how the energy is transported
from one place to other places in the flare region. The nature of the slow
evolution and long duration of the flare studies here provides us an excellent
opportunity to deduce a clear picture of how thermal plasmas evolve after
being heated by magnetic reconnection. Moreover, analyzing the sequential
evolution of the thermal sources along the PIL helps improve our understanding
of three-dimensional aspects of the flare process.

We derive the thermal properties of this flare based on imaging data from six
AIA EUV passbands, including 131Å (Fe XXI, $\sim$11 MK; Fe VIII, $\sim$0.4
MK), 94Å (Fe XVIII, $\sim$7.1 MK; Fe X, $\sim$1.1 MK), 335Å (Fe XVI, $\sim$2.5
MK), 211Å (Fe XIV, $\sim$2.0 MK), 193Å (Fe XII, $\sim$1.6 MK; Fe XXIV,
$\sim$17.8 MK), and 171Å (Fe IX, $\sim$0.6 MK) (O’Dwyer et al., 2010). We
adopt the method of Cheung et al. (2015), who use a sparse inversion code to
calculate the emission measure (EM) as a function of temperature from AIA
imaging data. This sparse inversion code is further updated by (Su et al.,
2018), who adjusts the parameters of the sparse code to better suppress
spurious high EM values at high temperatures. Thus, the new differential
emission measure (DEM, describing the amount of thermal plasma along the line
of sight (LOS) as a function of T) diagnostic derived from the same AIA data
is much more consistent with thermal X-ray observations.

In our calculation, we have re-binned the AIA images into a pixel size of
1.2$\arcsec$ (2$\times$2 rebinned) and used average intensities from two
adjacent frames at 24s time cadence (a rebinning of 2$\times$ in time) for a
better signal-to-noise ratio. The obtained EM is the line-of-sight integrated
measure per unit area across the images.

Here, we use the EM-weighted temperature $T_{\text{EM}}$ per pixel defined in
the following formula (Su et al., 2018) to construct the temperature map in
spatial domain:

$T_{\text{EM}}=\frac{\sum{(\text{DEM}(T)\cdot\Delta T\cdot
T)}}{\sum(\text{DEM}(T)\cdot\Delta T)}=\frac{\sum{(\text{EM}_{T}\cdot
T)}}{\sum{\text{EM}_{T}}}$ (1)

From the temperature map (Figure 3(a)-(d) and the associated animation), one
can clearly identify three episodes of the flare thermal evolution: In the
first episode, a hot region stood out in the southernmost part of the active
region, which corresponded to the group of sheared loops in Figure 2(a), later
on, a propagation of thermal sources proceeded sequentially in space along the
long arm of the $\Gamma$-shaped PIL during the second episode (Figure 3(b)),
in the last episode, a hot arched region was newly formed, straddling the
short arm of the $\Gamma$-shaped PIL (Figure 3(c) and (d)).

It is worth noting that, around 16:28 UT, there appeared two discrete thermal
components both in the temperature map and the intensity contour of (thermally
dominated) 6-25 keV RHESSI image (Figure 3(c)), these two components
simultaneously existed in separated locations. The thermal source with a
higher projected position results from a continual migration along the
extended curved PIL from the southeast to the northwest, while the one with a
lower projected position was located at the top of a newly formed post-flare
loop, as manifested by its apparent rise motion. The higher thermal source is
relatively hotter than that of the lower thermal source. The centroid
locations of these two components are separated by $\sim$ 70 $\arcsec$.

High-resolution HXR spectroscopy provides a powerful complement to imaging
observations. The RHESSI spectrometer consists of an array of nine segmented
germanium detectors (GeDs). Each detector is segmented into a thin front
segment, which records photons from 3 keV to 2.7 MeV, with a resolution (FWHM)
of 1 keV (at 100 keV), and a thick rear segment built to detect photons from
about 20 keV to 17 MeV, with a resolution of 3 keV (at 1 MeV, Wigger et al.,
2004). Because of the strong attenuation below $\sim$ 6 keV and the K-escape
events (The majority of the counts recorded below 6 keV is K-shell photon
escaped from the germanium detector bombarded with high energy photons), no
information can be gained about the incident photon spectrum below 6
keV(Phillips et al., 2006). Thus the energy fitting range is restricted above
6 keV. The X-ray spectrum here is generated using the combined RHESSI front
detectors #1,3,5,6,8,9 to balance the resolution and signal-to-noise ratio
(SNR). Among the excluded detectors, #2 and #7 show significantly worse energy
resolution than the other detectors, the photon spectrum recorded in #4
appears abnormal during this event. The nonsolar background spectrum is
selected during the neighboring RHESSI nighttime just before and/or just after
the flare of interest. In order to ensure a reliable data set of RHESSI, care
is taken to avoid the effect of attenuator state changes, satellite night
times, South Atlantic Anomaly (SAA), and other complexities like photon pileup
and decimation of data due to instrumental overflood. For the spatially
integrated spectra, we used the forward modeling method implemented by the
Object Spectral Executive (OSPEX; Schwartz et al., 2002). OSPEX uses an
assumed parametric form of the photon spectrum and finds parameter values that
provide the best fit in a $\chi^{2}$ value relating the observed background-
subtracted photon flux with the predicted photon flux computed by folding the
assumed incident photon spectrum through the spectrometer response matrix
(DRM).

The thermal model (single temperature, $f\\_vth$ in OSPEX) provides the plasma
temperature T [keV] and volumetric emission measure EM [$\mbox{cm}^{-3}$] of
the thermal source. The temperature and emission measure are free parameters
while the relative iron abundance is fixed by default at the coronal value in
the CHIANTI atomic database (Dere et al., 1997; Landi et al., 2013). The
direct evidence (HXR observation of footpoint source or emission above 20 keV)
for high-energy energetic particles is absent during the flare impulsive phase
(see Figure 2, accompanying animation, and Figure 4), which makes the
nonthermal component fitting not considered. Based on the thermal evolution of
the flare (see Figure 3(a)-(d)), the spatially integrated spectra around 13:42
and 15:04 UT are fitted with a single-temperature thermal spectrum
($f\\_vth$), while the spectrum around 16:28 UT is fitted with two isothermal
functions ($f\\_vth$+$f\\_vth$). The fitting result of these intervals are
presented in Figure 4(a)-(c). It is found that the X-ray spectrum around 16:28
UT is well fitted by two distinct thermal components: a hot component
($\sim$11.9 MK) and a superhot component ($\sim$29.9 MK), yielding reduced
$\chi^{2}$ value $\sim$ 0.99 (Figure 4(c)).

To identify precise locations of these two distinct thermal components, we
trace this flare evolution in the dual views of SDO AIA and STEREO-A EUVI with
120° separation angle (Figure 5(e) and the associated animation). The two
thermal components are directly observed in the hot AIA 131 (10 MK) and/or AIA
94 (6.4 MK) passbands (Figure 3(e) and (f)), but are absent in cool AIA and
EUVI passbands. Luckily, by the time they cool down to the STEREO EUVI 195Å
temperature range, then the associated post-flare loop tops can be visible.
The AIA 193Å filter has a similar response function to the STEREO/EUVI
counterpart (193Å $\rightarrow$ 195Å). We use a routine called
“scc_measure.pro” (Thompson, 2009; Zhou et al., 2017) to determine the 3D
structure from combined EUVI images from STEREO-A and SDO. The routine is a
widget-based application that allows the user interactively to identify the
same features in both images, then the 3D coordinates are calculated by the
triangulation method. Tops of their post-flare loop are marked in Figure 5(a)
and (b), separately. It is clear that there exist two distinct loop systems
(Figure 5(b)) producing these two thermal components, one is a sequence of
flare loops stretching along the PIL from the south to the north, the other is
a compact flare loop concentrating on a small region at the northmost part.

##  4 Discussion and Conclusion

Imaging and spectroscopic observations of this event show that the superhot
plasma ($\sim$30 MK) is distinct, both spectrally and spatially, from the
usual $\sim$ 10–20 MK plasma. The flare consists of two separate loop
structures at different locations. The longer loop system develops with
persistently higher temperatures, after the second CME; the shorter loop
system dominates mainly the SXR emission. In the case of Caspi & Lin (2010),
the centroids of the spatially distinct superhot and hot sources are separated
by 10$\arcsec$. In our observations, their projected distance is 70$\arcsec$.
Also, the three episodes of HXR enhancements as proxied by the time derivative
of the SXR in Fig. 1(a) are associated with two CMEs, which implies that the
whole event went through three successive flare processes separately, rather
than an extended heating process. Imaging observations in Fig. 2 corroborate
the general pattern of the coronal loop and footpoint evolution, consistently
confirming the three enhancements of HXR emission. The superhot component in
Fig. 3(b) and (c) is the extended structure originating from the second stage
evolution. The hot component in Fig. 3(c) and (d), however, originated from
the top of the post flare loop emerging in the last stage evolution.

To summarize the observations, the overall evolution of the two thermal
components, including their actual location of the creation site, and the
timing and relationship to the flare-energy release, is summarized as follows:

Accompanied with the first CME, the apparent sequential evolution along the
PIL is attributed to numerous episodes of similar magnetic reconnection
occurring successively along the long arm of the $\Gamma$-shaped PIL. Magnetic
reconnection commences in the southernmost part of the active region resulting
in a superhot region beneath. The thermal energy is then transported along the
magnetic field lines toward the footpoint of the magnetic loops via thermal
conduction, producing the flare ribbon and chromospheric evaporation toward
the loop top. The evaporated plasma is supposed to be further heated at the
loop-top region, contributing to the increase of the emission measure of the
hot plasma there. Later on, the other distinct loop system appears along the
short arm of the $\Gamma$-shaped PIL and produces the second CME. Initiated by
the magnetic reconnection, similar energy release and transport process occur
in this set of loop system. Compared with the first flare loop system, this
set of loop system is compact and its projection height is lower.

From a side view including the Earth perspective (e.g., Figure 2(c)), the
superhot component is located above the flare loop top in projection, similar
to the cases reported by Svestka & Poletto (1985); Caspi & Lin (2010). But
from the top view (e.g., Figure 5(b)), the two distinct thermal components are
associated with two separate loop structures with the longer one having higher
temperatures, which is reminiscent of the cases reported by Den & Somov
(1989); Nitta & Yaji (1997). It is not rare for homologous eruptions to
originate from different segments of the same extended PIL within a short time
interval (e.g., Shen et al., 2013; Liu et al., 2017). With different flare
loop systems cooling with different timescales and with the flare maximal
temperature correlated with the flare class (Caspi et al., 2014), it is
anticipated that flare plasmas in a set of homologous flares may be manifested
as two or more co-existing thermal components of different temperatures in
HXRs, which are most likely distinct in space, but in certain circumstances
may also appear co-spatial in projection (e.g. Sharykin et al., 2015).

In fact, this unique event has the longest impulsive phase of any M or X-class
flare in the present Hale cycle since 1995, based on a search of the GOES
database111http://sprg.ssl.berkeley.edu/$\sim$tohban/wiki/index.php/The
Slowest Flare. The slow and large-scale evolution provides an excellent
opportunity to perform an intricate structure of the thermal source locations
and study their evolutions. And, thanks to multiperspective observations, this
is the first time attempting to resolve the locations of the hot and superhot
sources in 3D, revealing the two sources coming from two different loop
systems instead of different altitudes of the same loop system.

The authors wish to express their special thanks to the referee for
suggestions and comments which led to the improvement of the paper. The
authors thank Säm Kruker, Wei, Liu, and Brian R. Dennis for constructive help
on RHESSI data analysis and the teams of _SDO_ , _RHESSI_ , _STEREO_ and
_GOES_ for excellent data set. This work is supported by the B-type Strategic
Priority Program XDB41000000 funded by the Chinese Academy of Sciences. Z.J.
is supported by NSFC grants 42004142 and LDSE201703 from Key Laboratory of
Lunar and Deep Space Exploration, CAS. R.L. and Y.W. are supported by NSFC
grants 41574165, 41761134088, 41774150, 11925302, and 41774178. J.S., M.D.,
and X.C. are supported by NSFC grants 11722325, 11733003, 11790303, 11790300,
and Jiangsu NSF grant BK20170011. L.L. is supported by NSFC grant 11803096.
J.C. acknowledges support by NSFC grants 41822404, 41731067, 41574170,
41531073.

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Figure 1:  (a) Temporal evolution of GOES soft X-ray flux (black), its time
derivative (red), temperature (gold), emission measure (dark green), and
RHESSI 6-12 keV count-rate (purple) of the observed M1.7 flare. According to
the three peaks in the time derivative of GOES SXR (red profile), three
episodes are marked by red segments in the lower-left of this panel; (b) The
height-time plots for the leading edge of two flare-associated CMEs, the
heights $R/R_{sun}$ (in solar radii with respect to the disk center) are
measured at the fastest segment of the leading edge, the CPA, which is defined
as the midangle of the two side edges of the CME in the sky plane, represents
the location of the CME. The Velocity (Vel) gives the linear speed of the
CMEs.

Figure 2: $Top:$ An overview of the evolution of the flare on July 17, 2012.
(a)-(d) Composite images of AIA 131Å (cyan) and 1600Å (red) and HMI line-of-
sight magnetic field (gray) at 13:42, 15:04, 16:28 and 16:36 UT, respectively.
The overplotted contours in yellow indicate the emission of 6-25 keV from
RHESSI at levels of 50% and 80% with its maximum marked by a plus symbol
(red). An animation of these panels is available starting on 17 July 2012,
12:00:40 until 17 July 2012, 18:58:40 UT. The video duration is 17s. $Bottom:$
Evolution of the flare footpoint ribbons. The sequentially brightening of the
flare ribbons are plotted with different colors from blue (the earliest time)
to red (the lasted time). The background gray-scale image is the radial
magnetic field of HMI magnetogram at 12:00 UT. (Animation of this figure is
available.)

Figure 3: $Top:$ panels (a)-(d), the thermal evolution of the flare at 13:42,
15:04, 16:28, and 16:36 UT, RHESSI hard X-ray (6-25 keV, black) sources shown
by the contours overlaid on the temperature map, The contour levels are 50%
and 80% of the peak flux; the plus symbols denote the derived centroid
locations of the two distinct thermal components (purple and red). An
animation of these panels is available starting on 17 July 2012, 12:00:20
until 17 July 2012, 16:59:32 UT. The video duration is 31 s. $Bottom:$ SDO/AIA
observations from each of the six coronal filters during 16:28 UT of the
17-Jul-2012 flare, panels (e)-(g) are reverse color images, the hot features
in (e) and (f) are identified by the red arrows which are absent in the other
AIA channels. (Animation of this figure is available.)

Figure 4: Panels (a)-(c): Photon flux spectra (black), model fit (hot
component: green; super-hot: purple; total model: red), nonsolar background
(grey), and normalized residuals during three peaks ($\sim$ 13:40, 15:05 and
16:28 UT), detectors 1F,3F,5F,6F,8F, and 9F are used

.

Figure 5: $Top:$ SDO/AIA 193Å and STEREO-A/EUVI 195Å images at about 17:40 UT
during the eruption. The plus symbols denote the post flare loop tops used for
3D triangulation. The black dotted dash lines depict the location of the hot
components, An animation of these panels is available starting on 17 July
2012, 12:00 until 17 July 2012, 18:59 UT. The video duration is 1m24s.
$Bottom:$ The dotted dash lines in panels (c) and (d) show the extended flare
loop system at 16:00 UT and the compact flare loop system at 16:50 UT; Panel
(e): Positions of the STEREO-A/B and Earth (SDO) in the ecliptic plane on 2012
July 17. The red dot on the Sun marks the flare source region, which appears
on the solar disk when viewed from STEREO-A, on the limb from SDO, and on the
backside of the Sun from STEREO-B. (Animation of this figure is available.)