# A Transient Coronal Sigmoid in Active Region NOAA 11909: Build-up Phase,
M-class Eruptive Flare, and Associated Fast Coronal Mass Ejection

Hema Kharayat  Bhuwan Joshi  Prabir K. Mitra  P. K. Manoharan  Christian
Monstein  Udaipur Solar Observatory, Physical Research Laboratory, Udaipur,
313 001, India  Radio Astronomy Centre, National Centre for Radio
Astrophysics, Tata Institute of Fundamental Research, Udhagamandalam (Ooty),
643 001, India  Arecibo Observatory, University of Central Florida, Puerto
Rico, USA.  Istituto Ricerche Solari Locarno (IRSOL), Via Patocchi 57, 6605
Locarno Monti, Switzerland

###### keywords:

Active Regions, Magnetic Fields; Coronal Mass Ejections, Low Coronal
Signatures; Flares, Pre-Flare Phenomena, Impulsive Phase, Relation to Magnetic
Field; Radio Bursts, Dynamic Spectrum, Type II, Type III

##  1 Introduction

S-Introduction

Solar flares are characterized by the sudden enhancement of localised
brightness in the solar atmosphere. During a flare, a tremendous amount of
energy, ranging as high as 1028\- 1032 erg, is released which manifests its
signatures in the entire electromagnetic spectrum i.e. $\gamma$-rays to radio
waves (e.g. see review by Shibata and Magara, 2011). It is now well understood
that a flare may be accompanied by the expulsion of large amount of plasma
from the corona into the interplanetary space, i.e. Coronal Mass Ejection
(CME). Contrary to such ‘eruptive flares’, a second class of flares known as
‘confined flares’, do not show any association with CMEs. The confined flares
mostly undergo evolution in compact loop networks without causing significant
reconfiguration of the large scale magnetic configuration (see e.g. Kushwaha
et al., 2014). On the other hand, the eruptive flares and associated CMEs
cause large-scale changes in the coronal magnetic structure. Further, Earth-
directed CMEs are known for their geo-effectiveness. It is, therefore,
essential to study the source region characteristics of CMEs for a better
understanding of space weather phenomena.

Observations show that the CME-productive active regions exhibit some
interesting features during the pre-eruption phase such as coronal sigmoids,
filaments (or prominences), filament channels, and extreme ultraviolet (EUV)
hot channels (see reviews by Gibson et al., 2006; Alexander, 2006; Chen, 2011;
Toriumi and Wang, 2019). Among these features, filaments are the most studied
structures associated with solar eruptions. Filaments are dark thread-like
structures which consist of cool, dense plasma material embedded in the hot
tenuous solar corona (Martin, 1998). These structures are traditionally
observed in ground based H$\alpha$ images. When observed above the solar limb,
these structures appear bright against dark background, hence, are termed as
prominences (Parenti, 2014). Filaments exist along the magnetic neutral line
or polarity inversion line with their legs rooted in the opposite polarity
magnetic flux regions. The narrow lanes between opposite polarity magnetic
fields where a filament can be formed, are called filament channels (Engvold,
1997). Space based observations have revealed important hot-plasma structures
that form in the active region corona, such as, sigmoids and hot channels.
Sigmoids are S (or inverse S) shaped structures which are observed in soft
X-ray (SXR; Manoharan et al., 1996; Rust and Kumar, 1996) and EUV emissions
(Cheng et al., 2014; Joshi et al., 2017a). With the advent of multi-channel
EUV imaging from Atmospheric Imaging Assembly (AIA; Lemen et al., 2012) on
board Solar Dynamics Observatory (SDO; Pesnell, Thompson, and Chamberlin,
2012), hot cohrent plasma structures (often termed as hot channels) have been
identified from several CME producing active regions (Cheng et al., 2013;
Patsourakos, Vourlidas, and Stenborg, 2013; Nindos et al., 2015). These
features essentially provide evidence of the magnetic flux rope: a set of
magnetic field lines winding around a common axis (Gibson and Fan, 2006).
Contemporary, multiwavelength studies suggest these hot channels to be the
earliest signature of a CME in the source region (Cheng et al., 2011; Joshi et
al., 2017a; James et al., 2017; Mitra and Joshi, 2019; Sahu et al., 2020).

The formation mechanisms of the flux ropes in the corona have been addressed
in many studies. It has been suggested that the presence of the coronal flux
ropes is intrinsically related to their emergence from subphotospheric layers
(Amari, Luciani, and Aly, 2004; Archontis et al., 2004). Further, it has been
recognised that a magnetic flux rope can be formed above the photosphere by
shearing motions, emergence or submergence of the photospheric flux, or by the
combinations of these processes; the above mechanisms essentially rely on the
flux cancellation at the photosphere (Kusano, 2005; Mackay and van
Ballegooijen, 2006; Green and Kliem, 2009; Green, Kliem, and Wallace, 2011;
Cheng et al., 2013). In view of the onset of an eruptive flare, the
observations suggest that the flux rope associated structures can be formed
either during a solar eruption or may exist well before it (Zhang, Cheng, and
Ding, 2012; Patsourakos, Vourlidas, and Stenborg, 2013; James et al., 2017;
Mitra, Joshi, and Prasad, 2020). The studies further indicates that the in-
situ build-up of the flux rope during or just before an eruptive event is
accomplished via magnetic reconnections (Cheng et al., 2011; Patsourakos,
Vourlidas, and Stenborg, 2013; Kumar et al., 2016). It is noteworthy that the
existance of the flux rope in CMEs have been evidenced by near-Earth
spacecraft measurements, in particular, by the in-situ detection of magnetic
clouds (Démoulin, Janvier, and Dasso, 2016; Wang et al., 2018; Syed Ibrahim et
al., 2019). Although the origin and basic structures of the flux ropes have
been recognised and studied in many previous studies, there is a lot more to
learn about the formation process of the flux ropes and temporal aspects of
their evolution with respect to the onset of CMEs in the source region.
Simultaneous multi-channel observations obtained from Helioseismic Magnetic
Imager (HMI; Schou et al., 2012) and AIA onboard SDO at high spatial and
temporal resolutions have provided us with a unique opportunity for such
investigations. The multi-wavelength observations ensure the examination of
various layers of the solar atmosphere right from the photosphere to the lower
corona. Once a flux rope formed in the corona, the dips in the helical
magnetic structure provide support for the filament material against gravity
(Kuperus and Raadu, 1974; Gibson and Fan, 2006, see also review by Gibson,
2018). However, the presence of dipped magnetic structure does not provide the
necessity for the filament. The existance of the filament channel is more
frequent than the filament itself, pointing toward a magnetic environment
which is able to support the filament even if the filament does not exists
itself. The filament is often overlaid by the sigmoidal structure (Gibson et
al., 2002). The relationship between the filament and the sigmoid have been
discussed in many previous studies (Gibson and Low, 2000; Pevtsov, 2002;
Régnier and Amari, 2004; Joshi et al., 2017a, see also review by Gibson,
2018). By using Yohkoh SXR data and Hα full disk observation, Pevtsov (2002)
have studied sigmoid-filament association for six active regions. He found a
cospatial association between filament and sigmoid, through which he concluded
that both features are related to same topological structure. Based on three-
dimensional non-linear force free field modelling, Régnier and Amari (2004)
have investigated the relationship between filament and sigmoid in the active
region 8151. They have addresed that both the filament and the sigmoid can be
described by long twisted flux tubes.

Sigmoids indicate the presence of sheared and weakly twisted coronal field
lines with twist of around 1 turn in the field. These sheared and twisted
magnetic field structures can store free magnetic energy (Gibson et al., 2006)
and have been regarded as the source region of powerful CMEs. Using SXR images
from the Soft X-ray Telescope (SXT) on the Yohkoh satellite, Hudson et al.
(1998) examined the coronal structures associated with the halo CMEs and found
a characteristic pattern of sigmoid-to-arcade development with the passage of
a CME. Using SXT data, Canfield, Hudson, and McKenzie (1999) classified active
regions according to their morphology (sigmoidal or non-sigmoidal) and nature
of activity (eruptive or non-eruptive). They have found that the sigmoidal
active regions are prone to be eruptive than the non-sigmoidal regions. In the
case study of a major geo-effective CME from acive region NOAA 12371, Joshi et
al. (2018) found CME to be associated with a coronal sigmoid that displayed an
intense emission from its core before the onset of the eruption. Various
studies have been performed in order to understand the formation and eruption
of the sigmoids (Sterling et al., 2000; Savcheva et al., 2012; Joshi et al.,
2017a). Joshi et al. (2017a) studied the formation of a coronal sigmoid in the
active region NOAA 11719 and found that the development of sigmoid structure
occurred through the successive interaction between two J-shaped bundles of
loops. Mitra et al. (2018) examined two major X-class solar flares from the
sigmoidal active region NOAA 12673 and found that the complex network of
$\delta$-sunspots in the active region resulted to the formation of the
sigmoid in the hot EUV channel. Their study also revealed the flux rope
structure by coronal magnetic field modeling. The multiwavelength case studies
of the build-up processes of the coronal sigmoids and early evolutionary
phases of CMEs from these regions are therefore extremely important toward
understanding the origin of large-scale solar eruptions and development of
methods for their forecast.

In view of the above scientific motivations, here we present a comprehensive
multi-wavelength analysis of the formation of coronal sigmoid and subsequent
disruption from the region that resulted into a fast CME ($\approx$ 1085 km
s-1) and an M-class two-ribbon flare (SOL2013-12-07T07:29). The sigmoidal flux
rope was developed over a magnetically weaker and dispersed bipolar active
region (AR) NOAA 11909 on 07 December 2013. Further, sigmoidal structure
disrupted just $\approx$ 20 min after its complete formation; hence, it can be
justifiably called a ‘transient’ sigmoid. The flux rope activates during the
precursor phase of the M1.2 flare during which the flux rope slowly rose.
Subsequently, the flux rope moved to the phase of eruptive expansion with the
onset of impulsive flare emission and exhibited important observational
signatures in EUV and radio wavelengths. The paper is organized as follows: in
Section method, we discuss a brief account of the observational data and
method used for the study. Multi-wavelength observations and results are
presented in Section results. We discuss and interpret the results in Section
discussion. A summary of the results is provided in Section conclusion.

##  2 Observational Data and Method

method

For the multi-wavelength study of M1.2 flare on 07 December 2013 and its
source active region (NOAA 11909), we analyze data from multiple ground and
space borne instruments which are briefly discussed below:

  1. 1.

Observations of the different solar atmospheric layers in (E)UV wavelengths
are obtained from the Atmospheric Imaging Assembly (AIA: Lemen et al., 2012)
on board Solar Dynamics Observatory (SDO: Pesnell, Thompson, and Chamberlin,
2012). SDO/AIA provides solar images in seven EUV channels: 94 Å, 131 Å, 171
Å, 193 Å, 211 Å, 304 Å, and 335 Å; two UV channels: 1600 Å and 1700 Å; and one
white light channel: 4500 Å. AIA provides 4096 $\times$ 4096 pixel images at a
high resolution of 0”.6 per pixel with a cadence of 12 sec.

  2. 2.

For the photospheric observation of the Sun, we have used the data provided by
the Helioseismic Magnetic Imager (HMI: Schou et al., 2012) of SDO. HMI has
spatial resolution of 0”.5 per pixel. We used white light continnum intensity
images and line-of-sight (LOS) magnetograms of the solar photosphere with a
temporal cadence of 45 sec.

  3. 3.

CME observations are analysed by the C2 and C3 white light coronagraphs of
Large Angle and Spectrometric Coronagraph (LASCO: Brueckner et al., 1995) on
board the Solar and Heliospheric Observatory (SOHO: Domingo, Fleck, and
Poland, 1995). The field-of-view (FOV) for C2 and C3 coronagraphs are 1.5–6 R⊙
and 3.7–30 R⊙, respectively.

  4. 4.

To study radio emission associated with this eruptive flare, we have analyzed
the data obtained from extended Compact Astronomical Low-cost Low-frequency
Instrument for Spectroscopy and Transportable
Observatory111http://www.e-callisto.org/ (e-CALLISTO; Benz, Monstein, and
Meyer, 2005). The radio observations presented in this study are taken from
the Ooty station of the e-CALLISTO network.

  5. 5.

We have used the H$\alpha$ observation with the Solar Flare Telescope (SFT;
Ichimoto et al., 1991) operated at National Astronomical Observatory of Japan
(NAOJ), Mitaka. SFT is the ground-based instrument of Solar Science
Observatory.

  6. 6.

Soft X-Ray (SXR) images of the Sun has been taken from the observations of
Solar X-ray Imager (SXI) aboard Geostationary Operational Environmental
Satellite (GOES; Bornmann et al., 1996). SXI provides the full-disk images of
the Sun in X-rays with a temporal cadence of one minute.

In this study, coronal magnetic field lines are modeled by employing the
potential field extrapolation method by solving a Green’s function (Seehafer,
1978) which is available in the SolarSoftware library encoded in the
Interactive Data Language (IDL). As the boundary condition for potential field
extrapolation, we used z-component of photospheric vector magnetogram of 06:15
UT from hmi.sharp_cea_720s series of HMI/SDO at a spatial resolution of 0”.5
pixel-1 and temporal cadence of 720s. Here, the HMI magnetogram has a
dimension of 812$\times$462 pixels which corresponds to a size of
$\approx$294$\times$160 Mm. Extrapolation has been performed up to a height of
$\approx$160 Mm from the photosphere. The extrapolated field lines are
visualised by using Visualization and Analysis Platform for Ocean, Atmosphere,
and Solar Researchers (VAPOR; Clyne et al., 2007) software.

##  3 Multiwavelength Observations and Results

results

###  3.1 Overview of Flaring Active Region NOAA 11909

S-event overview

On the day of the reported activity i.e. 07 December 2013, the AR NOAA 11909
was situated at the heliographic coordinate S16W49. This AR emerged from the
eastern limb of the Sun on 28 November 2013 as $\beta$-type sunspot and moved
to the far side of the Sun over the western limb on 10 December 2013 as
$\alpha$-type. During its lifetime, its magnetic configuration evolved as
$\beta$ \- $\beta\gamma$ \- $\beta$ \- $\alpha$\- type and produced three
C-class and only one M-class solar flares.

Figure 1.: Multi-wavelength view of AR NOAA 11909 on 07 December 2013 prior to
the event, (a) HMI white light image of AR showing the distribution of
sunspots. Small and large size sunspots are indicated by green and blue
arrows, respectively. (b) Cotemporal HMI LOS magnetogram showing the magnetic
field configuration of the AR 11909. Dashed yellow curve indicate the
approximate polarity inversion line (c) Mitaka H-alpha image showing the
presence of filaments F1 and F2 (marked by red arrows) at the studied AR.Fig1

Figure Fig1 shows the multi-wavelength view of AR NOAA 11909 on 07 December
2013 at 05:49 UT, prior to the onset of the M-class flare. Figures Fig1a and b
display the photospheric continnum intensity image and LOS magnetogram,
respectively. It is clear from these panels that the AR consists of several
small to medium size sunspots with dispersed magnetic field. Among these
sunspots, the largest one (indicated by blue arrow) form the main sunspot of
the leading group. Comparison of Figures Fig1a and b confirms that the leading
sunspot group largly consists of positive magnetic polarity while the trailing
part of the AR is of negative polarity. Notably, magnetic flux in the trailing
sunspot group is quite dispersed and composed of a couple of smaller sunspots
(shown by green arrows) along with small flux element of the plage region. The
positive and negative flux regions exhibit a clear seperation so that a
polarity inversion line (PIL) can be easily drawn which is delineated by a
dashed yellow curve in Figure Fig1b. The clear distinction of bipolar region
suggests this AR to be a $\beta$-type, prior to the reported activity.
H$\alpha$ image of AR shows the presence of two filaments: F1 and F2, lying
approximately along the PIL which are indicated by red arrows in Figure Fig1c.
Filament F1 lies in the northern part while filament F2 lies in the southern
part of the PIL.

Figure 2.: (a) AIA EUV image of the AR 11909 in 171 Å channel on 07 December
2013 at 06:24 UT processed by multi-scale gaussian normalisation technique.
HMI line-of-sight magnetogram contours are overplotted to show the
distribution of the photospheric flux. The positive and negative polarities of
the flux are shown by red and blue contours, respectively. Contour levels are
set as $\pm$150, $\pm$400, and $\pm$600 G. White and green arrows show the
closed and open loops, respectively. (b-c) Modeled coronal magnetic field
configuration based on the potential field extrapolation method. Different
sets of closed and open field lines are shown by different colors, from front
and side views, in panels (b) and (c), respectively.Fig2

In Figure Fig2a, we show AIA 171 Å image of the active region at 06:24 UT. For
a clear visualisation of coronal features, the image is processed by Multi-
Scale Gaussian Normalisation (MGN; Morgan and Druckmüller, 2014) technique,
which is based on localised normalisation of the data over a range of spatial
scales. In order to display the distribution of the photospheric LOS magnetic
field, we have overplotted HMI magnetogram as contours. Red and blue contours
are used to distinguish positive and negative polarity fluxes, respectively.
Two distinct sets of coronal loops can be observed from the AIA 171 Å image of
the AR (Figure Fig2a) that resemble a closed (white arrows) and an open (green
arrows) magnetic field configurations. Notably, the following sunspot group is
associated with closed field lines while closed as well as open field lines
originate from the leading sunspot group. To confirm the presence of open
field lines in the AR, we have extrapolated coronal magnetic field
configuration by the use of potential field extrapolation method (see Section
method). In Figures Fig2b and c, we display extrapolated coronal field lines
from front and side views, respectively. We noticed the presence of different
sets of closed coronal loops, connecting the positive flux region of leading
sunspot group with the negative flux of the following sunspot group. The AR
also contains open field lines which originate from the positive polarity
regions of the leading sunspot group (shown by blue and cyan color lines).

###  3.2 Build-up of Coronal Sigmoid in the Active Region

S-coronal sigmoid

Figure 3.: Series of selected AIA EUV 94 Å images showing the development of
coronal sigmoid. Two arrows in panel (a) indicate two sets of coronal loops.
Arrow in panel (b) marks the connectivity between two sets of loops. The
dashed curve in panel (e) indicates the sigmoidal structure in the AR 11909.
White box in panel (f) marks the region used to construct the AIA light curves
which are plotted in Figure Fig3b.Fig4

Figure 4.: An overview of the sigmoid in (a) AIA EUV 94 Å, (b) 131 Å, and (c)
soft X-ray data from GOES-15. Dashed curve in all the panels marks the
sigmoidal structure.Fig

Figure 5.: (a) GOES SXR flux in 1–8 Å and 0.5–4 Å channels on 07 December 2013
from 06:30 UT to 09:30 UT. The vertical lines are used to differentiate the
different phases of the M1.2 flare. (b) Cotemporal AIA light curves in 94 Å,
171 Å, 193 Å, 211 Å, 304 Å, and 1600 Å channels, normalized by their
corresponding peak fluxes. In order to get clear visualisation, AIA light
curves are further normalized by factor 3.5, 1.7, 2.7, 4.9, 5, and 2 for 94 Å,
171 Å, 193 Å, 211 Å, 304 Å, and 1600 Å, respectively. The region used to
construct the AIA light curves is shown by the white box in Figure Fig4f.Fig3

During $\approx$ 06:20 UT–07:00 UT, we observed signatures of the development
of a coronal sigmoid. In Figure Fig4, we present a series of AIA 94 Å images
which show sequential formation of sigmoid in the activity region. Much before
the onset of the flare ($\approx$ 1 hour), two bundles of coronal loops are
identified which are indicated by the arrows in Figure Fig4a. At this stage,
we cannot identify any connectivity between these two bundles of loops. From
$\approx$ 06:36 UT, a connection starts to develop between the loops (shown by
arrow in Figure Fig4b). From $\approx$ 06:36 UT–06:55 UT, the connecting loop
system undergoes through continuous evolution and expansion. As a result, we
observe the establishment of the sigmoid at $\approx$ 07:00 UT (shown by red
curve in Figure Fig4e) which provides an evidence for the presence of the
magnetic flux rope. A multi-wavelength view of the fully developed sigmoid is
presented in Figure Fig. Figures Figa and b display the sigmoid in EUV AIA 94
Å and 131 Å channels, respectively. Figure Figc shows the SXR image of the
sigmoid. For a clear identification of the sigmoid, we have marked the sigmoid
structure by a dotted curve in all of the panels. This sigmoidal flux rope
appeared till the begining of the M-class flare at $\approx$ 07:17 UT.
Afterward, the sigmoid disrupted as the eruption of the flux rope proceded
(described in Section S-fluxrope).

###  3.3 Evolutionary Phases of M1.2 Flare

phases Figure Fig3a depicts the temporal variation of the SXR flux in the 1–8
Å (shown by violet curve) and 0.5–4 Å (shown by red curve) wavelength bands of
the GOES during the time interval 06:30 UT–09:30 UT on 07 December 2013. The
evolution of the flare can be summarized by three phases: precursor
($\approx$07:00 UT–07:17 UT), impulsive ($\approx$07:17 UT–07:30 UT), and
gradual phase ($\approx$07:30 UT–08:55 UT). The precursor phase is
characterised by small-scale and gradual energy release during the slow rise
of the sigmoidal flux rope. For characterising the impulsive phase of the
flare, we have analyzed the temporal variation of the X-ray flux in the
available hardest X-ray energy band i.e. 3.0–25.0 keV which corresponds to the
0.5–4 Å wavelength band of GOES. A sharp increase in the X-ray flux at
$\approx$ 07:17 UT indicates the start of the impulsive phase of the flare. A
sudden enhancement in the intensity in most of the EUV channels at $\approx$
07:17 UT (discussed later in Figure Fig3b), also signifies the start of the
impulsive phase. The flare reached its peak intensity at $\approx$ 07:29 UT
after which X-ray flux decrease gradually, defining the gradual phase of the
flare. The evolutionary phases of the flare are described in Table Tab:1. The
temporal evolutions of (E)UV intensities in the 94 Å (6 MK), 171 Å (600,000
K), 193 Å (1 MK), 211 Å (2 MK), 304 Å (50,000 K), and 1600 Å (10,000 K)
wavelength bands of SDO/AIA are presented in Figure Fig3b. The region selected
for constructing the AIA light curves is shown by the white box in Figure
Fig4f. This region has been chosen because the flaring activity during the
studied period was occurred within this region. It is to be noted that the
intensity variation at the hot AIA 94 Å channel resembles the GOES soft X-ray
profiles with some time delay in their peak flux which indicate that, in spite
of different temperature sensitivities of the instruments, both emissions
originate from hot flaring plasma filled in the coronal loops. Interestingly,
time profiles of the intensity variations in other AIA channels (except 1600
Å) display two stage evolution. In the first stage, a sudden rise and fall of
the intensity is observed while in the second stage the intensity variation is
rather gradual. The intensity variation in AIA 1600 Å channel shows only
single stage evolution which resembles the first stage of other AIA channels,
viz. 171 Å, 193 Å, 211 Å, and 304 Å. Notably, the first stage of the intensity
variation in these channels is noticed earlier than the peak of GOES flux.

Table 1.: Summary of evolutionary phases of the M1.2 flare.Tab:1 Sr.No. | Phase | Duration | Remarks
---|---|---|---
1. | Build-up of |  $\approx$ 06:20 UT–07:00 UT | Two twisted coronal loop systems
| coronal sigmoid |  | are observed that sequentially join
|  |  | and form the coronal sigmoid.
2. | Precursor |  $\approx$ 07:00 UT–07:17 UT | The flux rope undergoes a slow rise;
|  |  | localized brightening suggesting
|  |  | small-scale magnetic reconnections.
3. | Impulsive |  $\approx$ 07:17 UT–07:30 UT | The flux rope undergoes eruptive
|  |  | expansion; development of two
|  |  | parallel ribbons; type III and type
|  |  | II radio bursts; “sigmoid-to-arcade”
|  |  | transformation begins.
4. | Gradual |  $\approx$ 07:30 UT–08:55 UT | A fast CME is first detected by
|  |  | LASCO C2 at $\approx$ 07:36 UT; highly
|  |  | structured post-flare loop arcade is
|  |  | formed; “sigmoid-to-arcade”
|  |  | transformation completes.

###  3.4 Photospheric Magnetic Field Variations

flux:variation

In order to understand the formation of sigmoid and triggering mechanism of
the eruptive flare, it is important to examine the photospheric magnetic field
changes in the AR. Such changes result in the storage of magnetic energy in
the corona and start to contribute much before the eruption. In Figure Fig5a
we provide HMI magnetogram on 06 December 2013 at 00:10 UT, one day prior to
the flaring activity. The region with significant photospheric changes is
marked by the black box in Figure Fig5a (see the animation associated with
Figure Fig5). To probe, further magnetic changes, we provide a series of
selected magnetograms of the marked subregion from 00:10 UT on 06 December
2013 to 05:10 UT on 07 December 2013. The green and blue boxes in Figures
Fig5b and Fig5j show the regions associated with the prominent changes in
negative and positive polarity fluxes, respectively. Comparing these two
panels, we can clearly observe a significant cancellation of the negative
polarity flux as well as decay of the positive polarity flux. A small element
of positive polarity is noticeable near the patch of negative polarity
(indicated by red arrow in Figure Fig5b). The negative polarity patch
undergoes continuous movement in northward direction and cancels out the
positive polarity element (marked by red arrow in Figure Fig5d). The negative
polarity patch continues its moving behaviour and disappears in the subsequent
intervel (marked by red arrows in Figure Fig5e–h). Few other moving magnetic
features of negative polarity are also noted near the PIL (indicated by yellow
and green arrows in Figures Fig5b–j) which also get cancelled out
subsequently.

In Figure Fig6a, we plot the positive and negative fluxes through the whole AR
(as shown in Figure Fig5a). For comparing magnetic flux variation with the
coronal energy release, we plot the GOES SXR flux in 1–8 and 0.5–4 Å wavelenth
channels in Figure Fig6b. We have found that the positive and negative
polarity magnetic fluxes show large variations before the onset of the flaring
activity. It decreases in a periodic manner. About five hour prior to the
event (at $\approx$ 02:00 UT on 07 December 2013), the positive and negative
fluxes start to decrease continuously. With the onset of the M1.2 flare
(indicated by vertical dashed line in Figures Fig6a and b), both the fluxes
maintain an almost steady level.

Figure 6.: (a) HMI magnetogram of the AR NOAA 11909 at 00:10 UT on 06 December
2013. Black box marks the region where we have observed the variations in the
magnetic flux. (b-j) Zoomed-in view of black box region showing the changes in
the magnetic flux. Green and blue boxes indicate the regions of large
variation in negative and positive polarity fluxes, respectively. Different
color arrows indicate different changing elements of magnetic flux near the
PIL. An animation of this figure is provided in the supplementary
materials.Fig5

Figure 7.: (a) The temporal change of the photospheric magnetic flux of
positive and negative polarities. (b) Cotemporal GOES SXR flux variation in
1–8 Å and 0.5–4 Å channels. The vertical dashed line in both the panels mark
the onset of the M1.2 flare.Fig6

###  3.5 Eruption of Sigmoidal Flux Rope

S-fluxrope

Examination of series of AIA 94 Å images suggests expansion of the sigmoidal
flux rope structure after $\approx$ 07:00 UT which subsequently proceeds with
its complete eruption. For a clear visualisation of the sigmoidal flux rope
eruption, a series of AIA 94 Å running difference images are presented in
Figure Fig9. It is observed that the eruption of the flux rope starts in the
precursor phase at $\approx$ 07:03 UT in the northwest direction with two
erupting fronts, which we call as the leading front (red arrows in Figures
Fig9c and d) and the following front (black arrows in Figures Fig9c and d).
The eruption of the leading front is clearly visible for an extended time
interval encompassing the precursor and impulsive phases of the flare.
However, the signatures of expanding following front are hard to distinguish
during the impulsive phase and subsequently diminishes. The leading edge of
the flux rope is shown by red arrows in Figure Fig9d. Coincidently, the
leading front of the erupting flux rope is associated with the “elbow” of the
sigmoidal structure. Unlike the leading edge, the eruption of the central
region of the sigmoidal structure is much less pronounced during the precursor
and main phases of the flare. Hence, to understand the kinematic evolution of
the flux rope at the source region we consider the height-time measurements of
the leading front.

Figure 8.: SDO/AIA 94 Å running difference images showing the eruption of the
flux rope between $\approx$ 06:54 UT and $\approx$ 07:20 UT. Two arrows in
panel (c) show the two erupting fronts of the sigmoid. Arrows in panel (d)
show the eruption of the sigmoid fronts.Fig9

In order to investigate the evolution of the sigmoidal flux rope eruption, we
plot a time-slice diagram for the duration 07:00 UT to 07:29 UT (Figure Fig10)
which include both the precursor and impulsive phases of the flare. For the
purpose, we have choosen a narrow slit S2S1 (Figure Fig10). Notably, the
eruption of the leading front of the flux rope was observed to follow this
path. The time-evolution of this slit is presented in Figure Fig10b. For
comparing the ascent of the flux rope with the evolutionary phases of the
flare, we overplot the GOES SXR flux in 1–8 Å channel (shown by red curve in
Figure Fig10b) which readily suggests that during the precursor phase of the
flare i.e. $\approx$ 07:00 UT–07:17 UT, the flux rope undergoes through a slow
rise with a linear speed $\approx$ 15 km s-1. After $\approx$ 07:17 UT i.e. as
the flare enters into the impulsive phase, the flux rope attains an
acceleration phase with sudden rise in its speed to $\approx$ 110 km s-1.

Figure 9.: (a) An AIA 94 Å image on 07 December 2013 at 07:16 UT with a
straight line S2S1 along which we have constructed a time-slice diagram. (b)
Time-slice diagram showing the eruption of the flux rope during 07:00 UT to
07:29 UT. Red curve shows the temporal variation of the GOES SXR flux for the
same time interval. In the precursor phase, we have observed a slow rise of
the flux rope with linear speed 15 km s-1. As the flare entered into the
impulsive phase, the flux rope eruption accelerated to high linear speed of
$\approx$ 110 km s-1.Fig10

###  3.6 Multi-wavelength Flare Evolution

S-flare evolution

In Figure Fig7, we present a series of a few selected AIA 94 Å images showing
the evolution of the AR during different phases of the M1.2 flare. A clear
erupting loop-like structure can be identified in Figure Fig7b. With the onset
of the impulsive phase (i.e. from $\approx$ 07:17 UT), we observed a set of
parallel ribbons (indicated as Rb1 and Rb2 in Figure Fig7d) which are situated
on both sides of the PIL. With the evolution of the flare, the brightness of
the ribbons enhances. Also, the separation between them increases. During the
gradual phase, we observed formation of highly structured post-flare arcade
encompassing the entire active region (Figures Fig7f–i).

Figure 10.: Series of selected SDO/AIA 94 Å images showing the evolution of
the flare during the main phase. Arrows in panel (d) and (e) indicate the two
ribbons of the M1.2 flare. An animation of this figure is provided in the
supplementary materials.Fig7

The AIA 304 Å images clearly show filaments in the precursor phase and their
subsequent evolution during the M-class flare. In Figure Fig8, we show a
series of AIA 304 Å images. To show the distribution of the positive and
negative polarities of the photospheric flux, HMI LOS magnetogram is
overplotted in Figures Fig8a and e. We noticed the existance of two filaments
along the PIL: F1 in the northern part (white arrows in Figure Fig8a) and F2
in the southern part (green arrows in Figure Fig8a) of the AR. Filament F1
coincides with the “elbow” of the sigmoidal structure implying a close spatial
association with the sigmoid (cf. Figures Fig7 and Fig8). Filament F2 is
situated far from the sigmoid, but within the AR. It is noticed that the
filament F1 activates and undergoes through eruption during the precursor
phase (white arrows in Figure Fig8b). However, no significant changes are
observed in the filament F2. The activation of the filament F1 starts at
$\approx$ 07:13 UT and it successfully erupts by $\approx$ 07:20 UT.
Afterward, we noticed the development of two-ribbon flare (Figure Fig8d). As
discussed earlier, the increasing seperation between the ribbons can clealy be
seen from Figure Fig8d–h (marked by white arrows in Figures Fig8d and f). The
post-flare loop arcade start to appear from $\approx$ 08:19 UT (Figure Fig8i).

Figure 11.: Series of selected SDO/AIA 304 Å images showing the evolution of
the flare during the main phase. White and green arrows in panel (a) show the
pre-existing filaments F1 and F2, respectively in the AR. The erupting
filament F1 is shown by white arrows in panel (b). Panels (d-h) and (i) show
two-ribbons (two white arrows in panels d and f) and post-flare loop arcade,
respectively in the AR. HMI line-of-sight magnetogram is overplotted in panels
(a) and (e). The positive and negative polarities are shown by yellow and blue
contours, respectively, with contour levels set as $\pm$150, $\pm$400,
$\pm$600 G. An animation of this figure is provided in the supplementary
materials.Fig8

###  3.7 Radio observations

Figure Fig11 displays the radio dynamic spectrum obtained from the CALLISTO
spectrograph within the frequency range of 50–300 MHz on 07 December 2013
during 07:17 UT–07:40 UT. For comparison, we overplot the GOES SXR flux in 1–8
Å wavelength channel within the same time axis (shown by white curve). From
Figure Fig11, we readily observe that during the impulsive phase of the flare,
a type III solar radio burst occurs at $\approx$ 07:24 UT. This type III burst
is followed by a split-band type II solar radio burst during $\approx$ 07:26
UT–07:29 UT. The fundamental band of type II burst is observed with starting
frequency of $\approx$ 173 MHz. Using the observations of fundamental band, we
have calculated the shock parameters using the Gopalswamy model (Gopalswamy et
al., 2013). The shock speed is estimated as $\approx$ 1143 km s-1 and the
average height of shock formation is $\approx$ 1.32 R⊙. The eruptive flare,
however, lacked solar proton event.

Figure 12.: Solar radio spectrum observed by CALLISTO spectrograph on 07
December 2013 from 07:17 UT to 07:40 UT within the frequency range of 50–300
MHz. Spetrograph shows a type III burst at $\approx$ 07:24 UT and a split-band
type II burst between $\approx$ 07:26 and $\approx$ 07:29 UT. White solid
curve over the spectrum shows the GOES SXR flux variation in 1–8 Å channel.
Two dash curves indicate the fundamental and harmonic bands of the type II
radio burst.Fig11

###  3.8 CME in the near-Sun region

According to SOHO/LASCO CME catalogue, the flux rope eruption is associated
with a halo CME which is detected by the LASCO C2 and C3 coronagraph (Figure
Fig12). The height-time plot corresponding to the CME within the LASCO FOV is
provided in Figure Fig13. The CME is first detected in the LASCO C2
coronagraph at 07:36 UT at a height of $\approx$ 2.76 R⊙ and is visible in C2
till 08:00 UT (the C2 data points are denoted by asterisk symbol (*) in
Figures Fig13a and b). The CME is observed till 10:42 UT by C3, when it
reached the height of 20.36 R⊙ (denoted by diamond symbols in Figures Fig13a
and b). By applying a linear fit to the height-time data points (Figure
Fig13a), we calculate the linear speed of the CME to be $\approx$ 1085 km s-1.
By using second order fit to the height-time data points, the acceleration of
the CME is estimated as $\approx$ -41.68 m s-2. Thus, we can infer that the
flux rope eruption led to a deccelerated high speed halo CME. The speed of the
CME is calculated as $\approx$ 855 km s-1 at a height of 20 R⊙.

Figure 13.: Running difference images derived from SOHO/LASCO C2 (panel a) and
C3 (panel b and c) coronagraph showing the propagation of halo CME originated
from AR 11909 during the M1.2 flare on 07 December 2013.Fig12

Figure 14.: Height-time plot of the CME on 07 December 2013 between 07:24 UT
and 10:42 UT with linear (panel a) and second order (panel b) fitting.
Asterisk (*) and diamond shaped symbols show the height of CME observed by the
C2 and C3 coronagraphs, respectively, of SOHO/LASCO. Linear fit in panel (a)
provides the linear speed of the CME to be $\approx$ 1085 km s-1 while the
second order fit in panel (b) shows that the CME is deccelerating with
acceleration of $\approx$ 41.68 m s-2.Fig13

##  4 Discussion

discussion

We study the formation of a transient coronal sigmoid and its subsequent
disruption which led to an M-class flare and fast halo CME. The activities
occurred in the AR 11909 on 07 December 2013. We summarize the chronology of
the event in Table 1.

We have observed the build-up and formation phase of a transient coronal
sigmoid, an evidence of magnetic flux rope, in hot EUV channel of 94 Å. The
AIA 94 Å observations cleary reveal that the formation process of coronal
sigmoid witnessed the joining of two twisted loop-like structures with an
overall enhancement in the brighteness which likely suggests that
reconnection-coupling of the two flux ropes resulted into the build-up of a
large sigmoidal flux rope. About $\approx$ 20 min after the complete formation
of the coronal sigmoid, the eruptive activities started from the region which
ultimately led to a two-ribbon flare and fast CME ($\approx$ 1085 km s-1). The
eruption of sigmoidal flux rope caused large-scale restructuring of the active
region corona and typical “sigmoid-to-arcade” transformation is observed.
Transient sigmoids imply those sigmoids which appear for a shorter time
duration and usually observed to form prior to an eruptive flare (see e.g.
Pevtsov, Canfield, and Zirin, 1996; Joshi et al., 2017a). To understand the
evolution of the sigmoidal flux rope in the AR, we have explored the coupling
between different lower atmospheric layers i.e photosphere, chromosphere, and
transition region. The observations of the chromosphere and transition region
in H$\alpha$ (Figure Fig1c) and AIA 304 Å channels (Figure Fig8) indicate the
existence of two filaments along the PIL. Filament F1 coincides with the
sigmoid. Due to the spatial correspondance of filament (F1) and sigmoid, we
can infer that both features correspond to the same flux rope. Notably,
filament F1 completely erupts with the eruptive-expansion of the sigmoidal
flux rope. The existance of the filament has been widely accepted as evidence
of the magnetic flux rope (e.g. Joshi et al., 2017a, 2018; Mitra et al., 2020;
Sahu et al., 2020). Through the analysis of the HMI magnetogram, we observed
the moving magnetic features along with magnetic field changes near the PIL.
These features are supposed to play important roles in the formation of flux
rope and its subsequent eruption (e.g. Joshi et al., 2011; Kushwaha et al.,
2014; Vasantharaju et al., 2019; Mitra, Joshi, and Prasad, 2020).

The photospheric observations of the AR also depicts the flux cancellation
near the PIL. As a consequence, we found a decrease in the net flux of the AR.
It is likely that the flux cancellation would lead to the tether-cutting
reconnection (Moore and Roumeliotis, 1992; Moore et al., 2001) among opposite
polarity foot points of the low-lying short loops that are rooted just across
the PIL thereby facilitating the development of the flux ropes. With the
progression of the build-up phase of the sigmoid, the flux rope acquires a
large amount of magnetic free enery, which is released during the magnetic
reconnection in the form of solar flares and CMEs (e.g. Mitra et al., 2018;
Mitra, Joshi, and Prasad, 2020). The AIA 94 Å images suggest the coupling of
two twisted loop-like structures during the build-up phase of the coronal
sigmoid which correlates with our interpretation of observed flux cancellation
process close to the PIL. Successful eruption of the flux rope is found to be
followed by a two-ribbon solar flare of M1.2 class. The increasing separation
between the flare ribbons is also noticed in the study. The flare ribbons
represent the foot-points of the coronal loops involved in the coronal
magnetic reconnection process (e.g. see, Veronig et al., 2006; Joshi et al.,
2007, 2009). With the progression of a solar flare, the coronal loops rooted
at the opposite polarity sides of the conjugate flare ribbons evolve from
‘stressed’ configuration to nearly potential state (Joshi et al., 2017b).

The appearance of the sigmoidal structure in the AR is considered as a
progenitor of CMEs in the source region (Canfield, Hudson, and Pevtsov, 2000).
The early evolution of the CMEs is associated with the activation of the flux
rope. In this regard, we would like to further emphasize the two-phase
evolution of the sigmoidal flux rope: slow rise and fast rise. The early
dynamical evolution of the flux rope is identified by slow rise phase with a
speed of $\approx$ 15 km s-1 which points toward the initiation phase of the
CME (Zhang et al., 2001; Joshi et al., 2016; Mitra and Joshi, 2019). The
evolution of the flux rope undergoes sudden transition from state of slow rise
to fast rise with the drastic increase of its eruption speed to $\approx$ 110
km s-1. The fast rise state essentially marks the impulsive acceleration phase
of the CME (Zhang et al., 2001). Our measurement of the flux rope speed during
its eruptive phase is in agreement with the speeds of plasmoids, reported
mainly during the limb events (e.g. Manoharan and Kundu, 2003; Joshi et al.,
2016). Notably, the slow and fast rise phases of the sigmoidal flux rope are
well correlated with the precursor and impulsive phases of the flare,
respectively. In essence, the formation of transient sigmoid and activation of
EUV flux rope support the scenario of the development of CME in the low corona
(see e.g. Manoharan et al., 2001; Cho et al., 2009; Joshi et al., 2018)

The radio dynamic spectrum shows the occurance of type III and type II radio
bursts during fast rise phase of the flux rope or impulsive phase of the
flare. Type III radio burst is observed at $\approx$ 07:24 UT and type II
radio burst is observed during $\approx$ 07:26 UT–07:29 UT. Type III radio
bursts are produced due to accelerated electrons propagating along the open
magnetic field lines (see review by Reid and Ratcliffe, 2014). The rapidly
accelerating flux rope structure would drive a coronal shock. The type II
radio burst in the frequency range of $\approx$ 50–300 MHz implies the fast
eruption of the CME-flux rope system in the lower coronal heights ($\approx$
1.32 R⊙). Considering the projected speed of erupting flux rope as $\approx$
110 km s-1 within the AIA 94 Å FOV (see Figure Fig10), the initiation time and
formation height of type II radio burst correlate with the imaging
observations. The CME is later detected on C2 coronagraph of LASCO at a height
of 2.76 R⊙ at 07:36 UT. Notably, the shock speed is found to be comparable to
the speed of the CME in the LASCO FOV.

##  5 Conclusion

conclusion

In this paper, we provide a multi-wavelength analysis of the build-up phase of
a sigmoidal flux rope in the corona and its subsequent “sigmoid-to-arcade”
transformation. Notably, the sigmoid studied here is of transient category and
its disruption started just $\approx$ 20 min after its complete formation. The
coronal sigmoid is associated with a simple $\beta$-type AR which posseses
relatively weaker and dispersed magnetic flux regions. Notably, both legs of
the sigmoid are rooted in the magnetically dispersed regions, devoid of
sunspots. The photospheric moving magnetic features and flux cancellation
events are observed during the build-up phase of the sigmoid as well as its
post-formation phase when the eruptive activities set-in. In particular, the
observed flux cancellation near the PIL has important bearing toward the
strengthening and subsequent activation of the magnetic flux rope by tether-
cutting magnetic reconnections. These small-scale, episodic reconnection
events will also contribute toward plasma heating which enhances the
temperature of the newly formed flux rope. At this pre-eruption stage, the
flux rope starts to become visible in SXR and hot EUV channels. The “sigmoid-
to-arcade” transformation accompanies a classical two-ribbon flare (of GOES
M1.2 class) and a fast CME ($\approx$ 1085 km s-1). The two-phase evolution of
the flux rope clearly bifurcates the precursor and impulsive phases of the
solar flare, which points toward the activation of flux rope by tether-cutting
process and subsequent large-scale reconnection driven by expanding CME in the
source region.

#### Acknowledgments

We would like to thank SDO team for their open data policy. SDO is NASA’s
mission under the Living with a Star (LWS) program. We also thank FHNW,
Institute for Data Science in Brugg/Windisch, Switzerland for hosting the
e-Callisto network. We also acknowledge National Astronomical Observatory of
Japan (NAOJ), Mitaka, Japan for providing the H$\alpha$ images of the Sun. We
are thankful to GOES for providing SXR data of the Sun. GOES is a joint
mission of NASA and the National Oceanic and Atmospheric Administration
(NOAA). We acknowledge VAPOR which is a product of the National Center for
Atmospheric Research’s Computational and Information Systems Lab, supported by
the U.S. National Science Foundation and by the Korea Institute of Science and
Technology Information. Authors would like to thank anonymous referee for
providing constructive comments and suggestions, which improved the
presentation and scientific content of the article.

Disclosure of Potential Conflict of Interest The authors declare that they
have no conflict of interest.

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