# Filament eruption from active region 13283 leading to fast halo-CME and
intense geomagnetic storm on 23 April 2023

P. Vemareddy  Indian Institute of Astrophysics, II Block, Koramangala,
Bengaluru-560 034, India

(Received September 13, 2023; Revised December, 10,, 2023; Accepted December,
14 2023)

###### Abstract

Using multi-instrument and multi-wavelength observations, we studied a CME
eruption that led to intense geomagnetic storm on 23 April 2023. The eruption
occurred on April 21 in solar active region 13283 near the disk-center. The AR
was in its decay stage, with fragmented polarities and a pre-existing long
filament channel, a few days before the eruption. The study of magnetic field
evolution suggests that the flux-rope (filament) has been built up by
monotonous helicity accumulation over several days, and further converging and
canceling fluxes lead to helicity injection change resulting in the unstable
nature of the magnetic flux-rope (MFR) and its further eruption. Importantly,
the CME morphology revealed that the MFR apex underwent a rotation of upto 56∘
in clockwise-direction owing to its positive helicity. The CME decelerates in
the LASCO-FOV and has a plane-of-sky velocity of 1226 km/s at 20 R⊙. In the
Heliospheric Imager FOV, the CME lateral expansion is tracked more than the
earthward motion. This implies that the arrival time estimation is difficult
to assess. The in-situ arrival of ICME shock was at 07:30 UT on April 23, and
a geomagnetic storm commenced at 08:30 UT. The flux rope fitting to the in-
situ magnetic field observations reveals that the MC flux rope orientation is
consistent with its near Sun orientation, which has a strong negative Bz-
component. The analysis of this study indicates that the near-Sun rotation of
the filament during its eruption to the CME is the key to the negative Bz-
component and consequently the intense geomagnetic storm.

##  1\. Introduction

Coronal mass ejections (CMEs) are large scale magnetized plasma structures
originating from the Sun. They propagate outward through interplanetary space,
affecting the planetary atmospheres (e.g., Webb & Howard 2012). Their
propagation toward Earth was established to be the cause of the most severe
geomagnetic disturbances on Earth e.g., Gosling et al. 1991; Gosling 1993;
Webb et al. 2000; Zhang et al. 2007. Therefore, it is of scientific and
technological interest to understand the complete picture of CMEs, including
their magnetic structure and the mechanisms involved in their origin from the
source regions, evolution, and propagation from the Sun to the Earth (Zhang et
al., 2021).

Historically, CMEs have been observed in the images of white-light
Coronagraphs onboard OSO7 (MacQueen et al., 1974), Skylab (Sheeley et al.,
1980), and Solar Maximum Mission (SMM; House et al. 1981). When a CME happens,
coronagraphs use Thomson-scattered light from free electrons in coronal and
heliospheric plasma to observe the outward flow of density structures coming
from the Sun. Following these early space coronagraphs, LASCO coronagraph
aboard SoHO, launched in 1995, is still continuously capturing white-light
observations of the Sun. From 2006, the coronagraph instruments onboard the
twin STEREO (Howard et al., 2008) spacecrafts started providing vantage point
white-light images in a much wider FOV, enabling one to track the CME
propagation up to 330 R ⊙. The CMEs in white-light images are often observed
in association with filament or prominence structures (e.g., Webb & Hundhausen
1987; Gopalswamy et al. 2003; Vemareddy et al. 2012b, 2017), X-ray sigmoids
(Canfield et al., 1999; Moore et al., 2001; Vasantharaju et al., 2019), and
EUV hot channels (e.g., Zhang et al. 2012; Cheng et al. 2013; Vemareddy &
Zhang 2014; Vemareddy et al. 2022) from the solar disk. Therefore, the white
light observations of the CMEs are frequently accompanied by eruptions from
the solar disk, e.g., by GONG H$\alpha$ telescope, the SOHO Extreme
Ultraviolet Imaging Telescope, the STEREO Extreme-UltraViolet Imager (EUVI),
and the recently launched SDO/AIA instrument. The in situ counterparts of the
CMEs are called interplanetary coronal mass ejections (ICMEs), whose
observations at 1 au, and behind are supplied by the spacecraft instruments
onboard, e.g., Voyagers, Ulysses, Helios, Wind, ACE, and STEREO etc.

The CMEs have a traditional three-part structure, which is usually understood
to be compressed plasma in front of a flux rope, a cavity, and a bright
filament or prominence surrounding the cavity (Illing & Hundhausen, 1985).
Therefore, the CME magnetic configuration is a magnetic flux rope (MFR) with a
helical-field wound around the central axis. After ejection, the CME largely
maintains its magnetic configuration or topology, as a twisted MFR and
therefore is able to continuously propagate outward through the heliosphere
interacting with the ambient solar wind. A statistical study by Vourlidas et
al. (2013) suggests that at least 40% of CMEs observed by space-borne
instruments have a clear MFR structure. According to many studies based on
satellite and ground-based data, this MFR structure in the outer corona is
thought to be an evolved form of the magnetic structure of filaments observed
in H$\alpha$ or sigmoidal structures in soft X-rays in the source regions.

Figure 1.— The CME eruption from the Sun occurred on April 21, 2023, at 18:00
UT. a) AIA 304 Å image of the Sun showing the location of the AR 13283 (inset)
as the source region of the CME eruption. b) AIA 131Åimages of the AR during
the eruption. c) HMI magnetic field observations of the AR 13283. The yellow
dashed curve traces the polarity inversion line between fragmented opposite
magnetic polarities. The long filament channel remains intact, and the
erupting feature is a manifestation of the sheared magnetic loops in the
vicinity of the PIL. d) LASCO/C3 running difference image delineates the
Earth-directed halo CME with bright leading edge and core.

Geoeffective CMEs are those that are directed toward the Earth and are called
Halo CMEs, therefore, the source region should be near the visible disk center
(Srivastava & Venkatakrishnan, 2004; Zhang et al., 2007). These front-side
halo CMEs associated with strong soft X-ray flares tend to be the most
geoeffective (Gopalswamy et al., 2007). After ejection from the inner corona,
the CMEs evolve with distinct kinematic, geometric, and magentic properties,
which determine their geoeffectiveness. The kinematic studies reveal that the
CMEs evolve with slow, fast, and gradual acceleration in the inner corona upto
0.1 au (Zhang et al., 2004). Its further propagation in the outer heliosphere,
is governed solely by the interaction of the CME and the ambient solar wind
via aerodynamic drag (Cargill, 2004; Vršnak et al., 2010) which causes the
faster CMEs decelerate and slower ones to accelerate, tending to equalise
their speeds with the speed of the ambient solar wind. Most importantly, the
CME is geoeffective when its interplanetary magnetic field is southward-
pointing, which depends on the orientation of the MFR with respect to the
ecliptic. Some CMEs are observed to rotate in the corona due to a variety of
mechanisms. Simulations showed the rotation is due to kink instability, and
the direction of the rotation depends on the handedness of the flux rope
magnetic field (Török & Kliem, 2003; Lynch et al., 2009). CMEs have also been
observed to deflect both in latitude and longitude (Isavnin et al., 2013).
Depending on the relative positions of the coronal features, the deflection
motion was frequently described as being toward the heliospheric current sheet
(Cremades & Bothmer, 2004) or away from coronal holes (Gopalswamy et al.,
2009). These geometric changes of the CMEs in the corona are the key for
explaining the CME impact on Earth. For example, Kay et al. (2017) studied the
trajectory of the series of CMEs launched from AR 11158, located near the disk
center, and found that only one CME encountered the Earth. The other missing
CMEs were found to deflect north up to 30∘ near the Sun.

Figure 2.— Emission signatures of the CME eruption from the AR 13283. (a)
Radio dynamic spectrum obtained from WAVES onboard STA. Arrows point to the
type II, III radio burst triggered by the eruption. (b) Radio dynamic spectrum
obtained by WIND/WAVES instrument located at the Lagrange L1 point. Arrows
refer to the type III bursts originating from the CME associated flare. (c)
Disk integrated GOES X-ray flux. The peak value of the flux refer to M1.7
class flare occurred from AR 13283. Vertical dotted lines refer to flare start
(17:44 UT), peak (18:12 UT) and end (18:44 UT) times respectively.

In the interest of space-weather impact, identifying the geomagnetic source
regions is very important to understand the characteristics of the magnetic
structure responsible for the CME eruption and its evolution in the sun-earth
line. In the 24th solar cycle, only three strong storms occurred with Dst
index $<-$175 nT, i.e., on 17 March 2018, 23 June 2015 and 26 August 2018
(Gopalswamy et al., 2022). In this study, we study the CME source region of
the intense storm occurred on April 23, 2023 with Dst Index -212 nT. The
source region is identified as active region 13283. An overview of the
observations is given in section 2. Results of the analysis on magnetic field
evolution in the AR, CME onset mechanism, CME propagation, and in situ
observations are described in Section 3. A summary of the results with a
discussion is provided in section 4.

Figure 3.— In-situ observations of the CME encounter the Earth’s atmosphere.
a) the Bz component of the magnetic field with time. The Bz started becoming
negative at around 07:30 UT on April 23. Blue, and magenta dotted vertical
lines correspond to shock (23/07:30 UT), and MC (24/01:00 UT) arrival times.
b) Dst index with time. The Dst index started decreasing to negative values
around 08:30 UT on April 23. It peaked at -212 nT at 05:30 UT on April 24,
2023, indicating an intense and long-duration geomagnetic storm. The storm’s
main phase (orange dashed arrow) persists for 22 hrs, followed by a long
recovery phase for about a couple of days. c) Histogram of Kp index values.
Red, yellow, and blue bars refer to index values 0-3, 3-6, 6-9, respectively.
The Kp index values reach up to eight during the storm’s main phase,
indicating an intense geomagnetic storm of G4 severity as per space weather
classification.  Figure 4.— Horizontal velocity of flux motions in the AR
13283. Background image is $B_{z}$ map with contours $\pm$150 G. The arrows
indicate the direction of motion, and their length is proportional to the
magnitude of a maximum of 0.7 km s-1. Large shear motions of the fluxes can be
noticed in the region enclosed by blue ovals. Also, fluxes emerge and move
apart during the evolution. In all panels, axis units are in pixels of 0.5”.

##  2\. Overview of Observations

We study the event of a large-scale CME eruption launched from the source AR
13283 on April 21, 2023. We use multi-spacecraft and multi-wavelength
observations to study magnetic characteristics of the source region, the
initiation of the CME from the source region, and its propagation toward the
Earth.

The source region observations are well covered by the instruments onboard the
Solar Dynamics Observatory (SDO). Atmospheric Imaging Assembly (Lemen et al.,
2012) observes the full disk of the Sun continuously in seven extreme
ultraviolet (EUV) channels. These channels correspond to the chromosphere in
the temperature range of around 20,000 K to the corona at 10 million K. The
cadence of the observations is 12s at a pixel scale of 0.6 arcsec. The
corresponding magnetic field measurements at the photosphere are obtained from
the Helioseismic and Magnetic Imager at a pixel size of 0.5 arcsec. Figure 1a
displays the full disk image of the Sun in AIA 304 Å waveband. The eruption’s
onset (CME launch) occurred at around 18:00 UT from AR 13283, located near the
disc center (S21∘W11∘) with a long filament channel present along its polarity
inversion line (PIL, yellow dashed curve). The observations of the AR in AIA
131 Å waveband and the corresponding magnetic field measurements at the
photosphere are shown in Figure 1b-c. The AR consists of fragmented polarities
without a major sunspot, which suggests that the AR is in the decaying phase
of its evolution. The solar feature of the explosion manifests from the
shearing and converging motion of the fluxes along the PIL. In fact, a long
filament channel preexisted along the PIL of the opposite magnetic fluxes. The
eruption feature is a filament as part of this multiply-threaded channel.
However, the left part of the long filament channel appeared unaffected during
the eruption. (Ref Figure 6).

The white-light observations of the CME are obtained from Large Angle
Spectrometric Coronagraph (LASCO; Brueckner et al., 1995) on board the Solar
and Heliospheric Observatory (SOHO) as well as from the Sun Earth Connection
Coronal and Heliospheric Investigation (SECCHI; Howard et al. 2008) onboard
the Solar-Terrestrial and Relational Observatory (STEREO; Kaiser et al. 2008).
LASCO provides white light images in the field-of-view (FOV) of 1.5-6 R⊙ (C2),
and 3-32 R⊙ (C3) to study the CME dynamics in the near-Sun corona.
Heliospheric Imagers onboard SECCHI image the sun in a much wider FOV (HI1:
15-90 R⊙ and HI2: 70-330 R⊙ ) and they are suited to understand the CME
propagation beyond the Earth. At the time of this CME launch, the STEREO-A
(STA) has a separation angle of 10∘ with respect to the Earth, so the CME is
viewed with a slightly similar morphology as with LASCO. Figure 1d presents
the running difference LASCO/C3 image showing the large scale halo CME from
the AR. The linear speed of the CME projected in the plane of sky is 1284 km
s-1which classifies the CME as fast. CMEs with speeds greater than 960 km/s
and widths greater than 66∘ are connected with type II radio bursts and are
highly geoeffective (Gopalswamy et al., 2001).

Figure 5.— Magnetic evolution in AR 13283. a) Net flux in positive (north) and
negative (south) magnetic polarities. Disk-integrated GOES X-ray (1.0–8.0 Å
passband) flux is also shown with the y-axis scale on the right. Grey bands
refer to the phase of the flux emergence (FE) and flux cancellations (FC) by
converging motions. Note that the M1.7 flare associated with the fast CME
occurred during the converging and canceling of flux motions. b) the helicity
injection rate ($dH/dt$) with time. Accumulated helicity, $H(t)$, is also
plotted with a y-axis scale on the right side.

The CME is associated with a GOES class M1.7 flare, as indicated by the disk-
integrated soft X-ray flux plotted in Figure 2c. The start, peak, and end
times are 17:44 UT, 18:12 UT, and 18:44 UT, respectively. The radio dynamic
spectrum obtained from the Radio and Plasma Wave Experiment (WAVES, Bougeret
et al. 1995) onboard the Wind spacecraft detects the launch of type III radio
bursts co-temporal with flare start and peak times (Figure 2b). In addition to
type III, the STA/WAVES dynamic spectrum detects type II radio bursts, which
arise from shocks driven by CME (Cane et al., 1987; Gopalswamy et al., 2019).
The shock front is seen to persist up to 20:30 UT, as evidenced in LASCO
white-light images, correspondingly, the type II emissions prolong.

The WIND satellite also records magnetic and plasma observations at the L1
point, which are obtained at a cadence of 92 seconds. The data demonstrate a
distinct in situ counterpart of the CME (ICME) on April 23, with a total field
intensity above 30 nT with a rotating magnetic field component (Ref Figure
11). Figure 3a depicts the observations of the z-component of the magnetic
field ($B_{z}$) with respect to time. The shock appears to have hit at 07:30
UT on April 23, with the $B_{z}$ component becoming increasingly negative
(southward $B_{z}$ component). With these observations of southward $B_{z}$,
the Disturbance-storm-time (Dst) index values111obtained from
https://wdc.kugi.kyoto-u.ac.jp/dst_realtime/202304/index.html, as shown in
Figure 3b, indicate the start of a geomagnetic storm. The Dst profile is well
correlated with the $B_{z}$ component associated with the shock-sheath, and
magnetic cloud regions. The storm peaked at -212 nT around 05:30 UT on April
24. This storm has a Kp index values of up to 8 and is classified as G4
severe. It is the first major storm in solar cycle 25 to cause stunning
northern lights (auroras) to extend to lower latitudes around the world.

Figure 6.— EUV images during the onset of the eruption. a) AIA 304 Å Image
showing the filament section being on rise motion. Flare ribbons are formed on
each side of the PIL. b-c) AIA 171 and 193 Å images show the erupting
filament. In all panels, contours of Bz ($\pm$150 G) are overlaid. These
images reveal the filament threads wound in right-hand direction.  Figure 7.—
Onset of the eruption from AR 13283. left column: AIA 304 Å difference images
showing the rising filament section at different times. A cusp shape is formed
underneath the rising filament after 18:02 UT. middle column: AIA 171
Ådifference images showing the rising filament and the oval shape disturbance
due to rising filament as a blob-like structure. right column: AIA 193 Å
difference images show the oval-shaped disturbance surrounding the rising
filament environment. The eruption scenario is illustrated in the panels, with
the orange color curve representing the erupting filament and the overlying
loops undergoing reconnection as pink curves. All panels are overlaid with Bz
contours at $\pm$150 G. To better understand the eruption onset, the image
sequence from 17:51-18:04 UT is attached as a video. (An animation of this
figure is available.)  Figure 8.— GCS fit to CME morphology observed from COR2
and LASCO. a-b) CME morphology observed from C2 and COR2, at 18:36 UT and
18:38 UT respectively. MFR axis orientation is indicated as seen in the source
region (orange) and LASCO (green) c-d) CME morphology observed from C3 and
COR2, at 20:36 UT and 20:38 UT respectively. In all the panels, the best-
fitted wireframe of the MFR in the GCS model is overplotted. Since the STA and
SoHO are separated at 10∘, the CME morphology is similar in both views.
Importantly, the MFR oriented from north-east to south-west direction with a
tilt angle of -20∘ at 18:36 UT and -26∘ at 20:36 UT. Compared to source region
PIL, the filament apex appears to rotated by about 56∘ during its eruption to
CME within 10 R⊙ Figure 9.— CME propagation from low corona into the
heliosphere. a-d) Combined running difference images of STA/COR2 (2.5-15 R⊙),
STA/HI1 and STA/HI2 during April 21–23, 2023. The yellow arrow points to the
leading edge of the CME at different epochs of its propagation. Note that
since the CME is propagating towards the STA spacecraft, the propagation in
the images mostly represents the lateral expansion of the CME. The sequence of
images displayed in each panel is attached as a video stream. (An animation of
this figure is available).

##  3\. Results

###  3.1. Magnetic Evolution of the AR

To study the magnetic field evolution in the AR, we obtain time series (every
12 minutes) of line-of-site magnetic field observations from the HMI. The
magnetic field measurements at each instant of time are projected to the disk
center by the cylindrical equal area projection method to minimize the
projection effects. Then we choose the area of AR extent such that the foot
points of the loops connected to the filament channel are covered over the
evolution time under study. From these magnetograms, we derive the horizontal
velocity of flux motions by using the Differential Affine Velocity Estimator
(DAVE, Schuck 2005) method. The velocity field is displayed in Figure 4a,
where horizontal velocities (arrows) are plotted on Bz images. Contours of 150
G are also overlaid to identify the boundaries of individual polarity regions.
The velocity field is spread up to a maximum value of 0.8 km s-1indicating
that different features move at different velocities. We found that emerging
flux regions moved apart after their emergence.

The time evolution of the net magnetic flux is plotted in Figure 5a. The net
flux in the north (south) polarity is around $6\times 10^{21}$ Mx ($8\times
10^{21}$ Mx). These fluxes are half of the values typically noticed in the
growing ARs consisting of large-scale sunspot polarities (Vemareddy, 2019;
Vemareddy & Démoulin, 2017). The net flux profile increases during the first
half of April 20 corresponding to the flux emergence, and 12 hours before the
eruption, the net flux decreases by 1$\times$1021 Mx in each polarity. From
the velocity and magnetic fields, we also calculated the helicity injection
rate ($dH/dt$) to evaluate the net helicity coming from the flux motions
(Démoulin & Berger, 2003; Vemareddy et al., 2012a; Vemareddy & Démoulin,
2017). The $dH/dt$ and $H$ with time are plotted in Figure 5b. The $dH/dt$
profile evolves with positive values around a mean value of $3.5\times
10^{36}$ Mx${}^{2}\,s^{-1}$ till April 21, after which it becomes marginally
negative and then oscillates around zero values. The flux emergence phase is
seen to pump additional helicity flux above the mean value. A positive sign of
helicity refers to the right-hand helicity of the magnetic structure, which is
also in agreement with the windings of the filament threads (Figure 6). During
the four days of evolution, the flux motions accumulate a coronal helicity of
$0.9\times 10^{42}\,Mx^{2}$. It is worth pointing out that the filament
channel was present well before April 18, therefore the magnetic structure
exists with pre-accumulated helicity in the corona. Accounting for the net
average flux in the AR as $7\times 10^{21}$ Mx, the average twist of the flux
tube ($H/\Phi^{2}$) could be $0.02-0.1$ turns, which is in the range of
several erupting ARs (Vemareddy, 2019). In conclusion, the large shear and
converging motions build the twisted flux about the PIL, which was set to
erupt during the phase when flux motions started injecting weak helicity flux
with alternate polarity into the AR corona.

###  3.2. Initiation and onset of the CME eruption

The magnetic evolution in the AR is favorable for the eruption of twisted
fields. Before the four days of the eruption, a long filament channel already
existed, lying above the PIL. Its one leg lies in the leading negative
polarity, and the other lies on the extreme east edge of the following
positive polarity. Figure 6 displays the EUV images showing the filament and
the surrounding loops. These images reveal the orientation of the filament
axis and the PIL, which are found in a south-east to north-west direction at a
tilt angle of 30∘. The filament threads, which are wound in the right-hand
direction, are clearly visible along the filament channel. The winding
direction reflects the right-handed helicalness or positive helicity of the
magnetic field, in agreement with the sign of the helicity injection
estimates. The overall shape of the filament is a forward S-shape, which
indicates positive helicity and is compatible with the hemispheric rule of
sigmoids being in the south hemisphere.

About two hours before the main eruption, the filament channel shows signs of
eruption, like brightening along its length, threads being reorganized, and
the surrounding loop structures being heightened. The initiation of the
eruption started at around 17:42 UT, with the sub-filament section along the
PIL starting to rise. This means that the large-scale filament is a
combination of sub-twisted loop systems originating from different locations
along the PIL. With this rise, the GOES X-ray flux also increases, commencing
the flare at 17:44 UT. Figure 5 delineates the rising filament section and the
two flare ribbons at 17:56 UT. In order to visualize the filament and the
surrounding loop system, we analyze running difference images, as displayed in
Figure 7. AIA 304 Å difference images show clearly the rising filament section
with time and its disappearance after 18:03 UT, whereas AIA 171 and 193
difference images give the impression of a blob-like structure expanding,
which forms an oval shape in exactly the same direction as the PIL, i.e.,
south-east to north-west. Further rise motion brought the filament into the
phase of the onset of its eruption, and this was clearly associated with the
formation of the cusp-shaped loops underneath the expanding filament. The
start of the filament rise motion at 17:42 UT and its onset at 18:03 UT fall
into the scenario of tether-cutting reconnection (Moore et al., 2001). The
reconnection of the surrounding less sheared loops adds more field to the
filament, which is regarded as the MFR, and the main phase of the eruption
occurs as a result of the run-away tether cutting reconnection, which is
linked to strong flare ribbons at the peak flare time. This eruption scenario
is illustrated in the right panels of Figure 7, with the orange color curve
representing the erupting filament and the overlying loops undergoing
reconnection as pink curves. With the benefit of vantage point observations,
such eruptions are being investigated more thoroughly, particularly the
kinematics of slow and fast rising motion (Vemareddy et al., 2012b, 2022).

###  3.3. Near sun rotation of the MFR

To determine the underlying large-scale structure and orientation of the MFR,
we exploited the plane of sky-projected coronagraph images from different
vantage points. For this purpose, we employed the Graduated Cylindrical Shell
(GCS; Thernisien et al. 2009) model. In this model, the large-scale structure
of the MFR is approximated either by the conical legs or the curved (tubular)
fronts. The model is based on height-aspect ratio ($\kappa$), half-angle
($\alpha$), tilt angle ($\gamma$), latitude ($\theta$) and longitude ($\phi$).
The underlying assumption of this model is that the pre-eruptive magnetic
configuration of the source AR constrains the magnetic orientation of an
erupting MFR. This is based on an analysis of several CMEs and the magnetic
configurations of their source regions (Cremades & Bothmer, 2004), which is
incorrect in a number of observed cases.

In Figures 8, we displayed the difference images of STA/COR2, LASCO/C2, and
C3, obtained at two different epochs. The wire frame (dotted green) from the
GCS fit is overlaid in these panels. Since SoHO and STA spacecrafts are
separated by 10∘, the morphology of the CME is observed to be similar from the
two different vantage points. The GCS fit to the CME morphology is implemented
with a GUI procedure rtsccguicloud.pro available in the STEREO tree of SSWIDL
software (Freeland & Handy, 1998). While fitting the GCS model, the tilt angle
of the erupting feature and the PIL orientation in the source region are
generally considered because the erupting feature (filament or prominence) is
assumed to be the MFR. However, we notice that the PIL orientation (south-east
to north-west) in the source region is quite different from the CME
morphology, which is in a north-east to south-west direction. It appears that
the MFR (erupting feature) from the source region has undergone a rotation of
about 60∘. Therefore, we use this information to set the tilt angle by
visually matching the observed CME morphology. From this fit, the tilt angle
is found to vary from -20∘ to -25∘ while the CME evolves from 18:36 UT to
20:36 UT. Also, the corresponding heights are noted as 10 R⊙ and 27 R⊙
respectively. From these two heights in time, we found that the CME velocity
in 3D space exceeds 1500 km s-1.

Figure 10.— CME kinematics in the plane-of-sky observations. a) A time-
elongation map (J-map) constructed from the slices at a position angle of
260∘. Red dots are the trace of the leading edge of the propagating CME. b)
Height-time data of the CME. The red curve is a second order fit to the LASCO
data (“$\diamond$” symbol), and the blue curve is a second order fit to J-map
data (“+” symbol). The Sun-Earth distance is shown with horizontal dotted
line, and vertical grey line represents the time (23/02:50 UT) of Sun-Earth
distance passage in the plane-of-sky.  Figure 11.— In-situ observations of the
CME (ICME) from WIND spacecraft. Panels from top to bottom show plots of
magnetic field strength ($B_{\rm tot}$) and its components ($B_{x}$, $B_{y}$,
$B_{z}$), and solar wind speed ($V_{\rm sw}$) as a function of time,
respectively. The blue dashed vertical line marks the arrival time of shock
(23/07:30 UT) at the spacecraft, and the orange vertical lines correspond to
the MC start and end times. The blue and red curves are the fitting results of
the Gold-Hoyle and Lundquist models, respectively. The fitting yields MC
orientations of -5∘ and -15∘ latitude respectively, indicating a mostly
southward Bz component.

###  3.4. Propagation of the CME

After the CME’s first emergence in LASCO C2 fov at 18:12 UT, its propagation
in the C3 and COR2 is very clearly observed in the running difference images.
Because of line-of-sight propagation, the CME motion in these images only
represents lateral expansion. The CME further propagation is tracked in the
wide-angle observations of the HI1 and HI2 of STA. The running different
images of individual instruments are combined to make a composite running
difference image at each instant of time so that the CME’s heliospheric
propagation is visualized continuously while it transits from one fov to the
other upto 370 R⊙. In Figure 9, we display these composite running images at
four epochs. The bright shock front is evidently noticed, along with the CME
leading edge in HI2. However, it is difficult to identify the precise
boundaries of the CME core and leading edge. The morphology is not spherically
symmetric rather a cone shape inclined roughly about 25∘ to the ecliptic
plane, consistent with the near Sun CME orientation.

The event is also cataloged in the HELCATS list222https://www.helcats-
fp7.eu/catalogues/event_page.html?id=HCME_A__20230421_02, providing useful
kinematic information from the analysis of time-elongation map (J-map)
constructed from combined difference images (Sheeley et al., 1999; Davies et
al., 2009) as displayed in Figure 10a. The time-elongation, in terms of
height, is plotted in Figure 10b. Also, we plot the height-time data obtained
from the LASCO CME
catalogue333https://cdaw.gsfc.nasa.gov/CME_list/UNIVERSAL_ver1/2023_04/univ2023_04.html.
A second order polynomial function best fits these observed height-time data
of CME propagation. Starting from a velocity of 1459 km/s, the CME decelerates
at a rate of 14 m s-2 in LASCO fov. At a height of 20 R⊙, the velocity of the
CME is 1226 km/s. In the HI fov, the travel time of 1 au distance corresponds
to 33 hours (3:00 UT on April 23) from CME onset at 18:00 UT on April 21. A
difference of 4.5 hours compared to the observed arrival time of the ICME
shock is due to projection effects. This emphasizes the difficulty of tracking
the earthward propagation of the CME in this case.

Since the CME launched at 20:30 UT on April 21 at the Sun, the time of arrival
of the shock front is 35 hours. As a geomagnetic storm commences at the same
time as a shock front’s arrival, it is important to assess the arrival time of
the CME frontal structure, which is ICME shock front. We use an initial CME
speed at 20 R⊙, in the drag based model. Also, the drag coefficient
$\gamma=0.15\times 10^{-7}$ km-1 and a background solar wind speed of 400 km/s
are used. For the observed 1226 km/s speed, the transit time exactly matches
with the arrival of the ICME sheath. Any higher value of drag coefficient may
require a higher value of CME speed to match with its observed in situ
arrival.

###  3.5. In situ magnetic field observations

The WIND observations of the in-situ magnetic field are plotted in Figure 11.
These observations are one minute averaged in the Geocentric Solar Ecliptic
(GSE) coordinate system. The shock front encounters the spacecraft at 07:30 UT
on April 23. Following the ICME shock-sheath region, the MC passes, which is
identified by the systematic variation of the magnetic field components,
indicating a strong magnetic field (30 nT) associated with an MFR-like
structure. The Bz component is rotating, with its sign changing from negative
to positive. Also, the solar wind velocity increases from the background
values of around 360 km s-1. Based on these observations, we identified the
time intervals for the MC as 01:00-15:00 UT on April 24.

The MC signatures are often modelled by an MFR. We employ the Lundquist (1950)
and Gold & Hoyle (1960) (GH) models of cylindrically symmetric linear force-
free magnetic fields. Because the MC grows during transit, the fitting
technique takes this into consideration (Vemareddy et al., 2016). The magnetic
field profile along the spacecraft’s observational path in the Lundquist model
is determined by the orientation of the flux rope axis, i.e., elevation and
azimuth angle ($\theta$ and $\phi$), the closest approach of the observational
path p, the flux rope diameter D, the helicity sign H, and the field strength
$B_{0}$ at the flux rope’s centre (Lepping et al., 1990). In the GH model, the
twist parameter $T_{0}$ is added to these parameters. The helicity sign is
decided on a trial basis; in this case, it is positive and consistent with the
helicity sign of the source region, and the rest of the parameters are
determined when the fitting converges in the least squares sense. As a measure
of goodness of the fit, we compute the root-mean-square deviation
($rms=\sqrt{\sum_{i=1,N}\left(\mathbf{B}_{o}(t_{i})-\mathbf{B}_{m}(t_{i})\right)^{2}/N}$)
between the observed magnetic field ($\mathbf{B}_{o}(t_{i})$) and the modeled
field ($\mathbf{B}_{m}(t_{i})$) (Marubashi & Lepping, 2007). In terms of rms
value, the GH fit (rms=6.76) is a better model than the Lundquist model
(rms=9.49) as compared to the observations in Figure 11. Based on these fits,
the MC axis is oriented -15∘ in latitude and roughly 260∘ in longitude. These
results indicate that the MC magnetic fields had a significant southward Bz
component to launch the geomagnetic storm. The flux rope has a positive sign
of magnetic helicity which is inline with the source region helicity sign of
the erupting part of the filament as seen in EUV images (Figure 6). The twist
of the magnetic field in the GH flux rope model is 1.67 turns/au. The MC has a
diameter of 0.17 au. While passing the spacecraft, the radius of the MC varies
with time: $r(t)=r_{0}(1+Et)$ where E is the expansion rate determined to be
0.09 per day from the MC fitting analysis.

Table 1Chronology of the CME eruption from AR 13283 21/17:44 UT | CME initiation & flare start time
---|---
21/18:03 UT | Onset of CME eruption
21/18:12 UT | peak time of M1.7 flare
21/20:30 UT | CME LE at a height 20 R⊙
23/07:30 UT | ICME shock arrival
24/01:00-15:00 UT | MC passage
23/08:30 UT | geomagnetic storm commencement
24/05:30 UT | geomagnetic storm peaked at 212 nT

##  4\. Summary and Discussion

We investigated the CME eruption that occurred on April 21, 2023 which caused
a severe geomagnetic storm on April 23, 2023. The eruption originated from
solar AR 13283 near the disk center. The AR was in its decay stage, with
fragmented polarities and a pre-existing long filament channel, a few days
before the eruption. The polarities have undergone large-scale shear and
converging motions, so that the flux cancellation leads to the buildup of
twisted flux threads along the main PIL. The magnetic helicity is positive, so
the chirality of the filament threads is right-handed. Importantly, the
helicity injection changed from a positive to a negative sign early on April
21. Because the flux rope (filament) has been built up by monotonous helicity
accumulation over several days (Amari et al., 2003; Vemareddy & Mishra, 2015),
the further converging and canceling fluxes lead to helicity injection change
and the unstable nature of the MFR and its further eruption.

The filament that is initiated to erupt is a part of the long, multiply-
threaded filament channel. Along with the increase in GOES X-ray flux, the
filament set to rise at 17:42 UT and eventually erupted at around 18:03 UT
(see Table 1). Importantly, the CME morphology revealed that the MFR is
oriented in the north-east to south-west direction (-26∘ with respect to the
equator), which differs by 56∘with respect to the filament orientation, which
is in the south-east to north-west direction (30∘ with respect to the
equator). These observations imply that the filament apex rotates in a
clockwise direction owing to the right-hand twist of the magnetic field (Lynch
et al., 2009; Green et al., 2011).

The CME decelerates in the LASCO fov and has a velocity of 1226 km/s at 20 R⊙.
In the HI fov, the CME lateral expansion is tracked more than the earthward
motion, then the arrival time estimation is difficult to assess. However, the
drag-based model gives correct estimates of the arrival time of the ICME
sheath with the initial CME speed at 20 R⊙. The in situ arrival of the ICME
shock was at 07:30 UT on April 23, and then a geomagnetic storm commenced at
08:30 UT. The flux rope fitting to the in situ magnetic field observations
reveal that the MC flux rope orientation is consistent with its near Sun
orientation within 10 R⊙, which has a significant negative Bz-component. The
analysis of this study indicates that the near-Sun rotation of the filament
during its eruption to the CME is the key to the negative $B_{z}$-component
and consequently the intense geomagnetic storm.

SDO is a mission of NASA’s Living With a Star Program; STEREO is the third
mission in NASA’s Solar Terrestrial Probes program; and SOHO is a mission of
international cooperation between the ESA and NASA. We acknowledge the use of
NASA/GSFC’s Space Physics Data Facility’s OMNIWeb (or CDAWeb or ftp) service.
The CME catalog used in this study is generated and maintained by the Center
for Solar Physics and Space Weather, the Catholic University of America, in
cooperation with the Naval Research Laboratory and NASA. The author is
grateful to the anonymous referee for a detailed list of comments and
suggestions.

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