# The Merging of a Coronal Dimming and the Southern Polar Coronal Hole

Nawin Ngampoopun University College London, Mullard Space Science Laboratory,
Holmbury St. Mary, Dorking, Surrey, RH5 6NT, UK  David M. Long Astrophysics
Research Centre, Queen’s University Belfast, University Road, Belfast,
Northern Ireland, BT7 1NN, UK  University College London, Mullard Space
Science Laboratory, Holmbury St. Mary, Dorking, Surrey, RH5 6NT, UK  Deborah
Baker University College London, Mullard Space Science Laboratory, Holmbury
St. Mary, Dorking, Surrey, RH5 6NT, UK  Lucie M. Green University College
London, Mullard Space Science Laboratory, Holmbury St. Mary, Dorking, Surrey,
RH5 6NT, UK  Stephanie L. Yardley Department of Meteorology, University of
Reading, Reading, UK  University College London, Mullard Space Science
Laboratory, Holmbury St. Mary, Dorking, Surrey, RH5 6NT, UK  Donostia
International Physics Center (DIPC), Paseo Manuel de Lardizabal 4, 20018, San
Sebastián, Spain  Alexander W. James University College London, Mullard Space
Science Laboratory, Holmbury St. Mary, Dorking, Surrey, RH5 6NT, UK  European
Space Agency, European Space Astronomy Centre, Spain  Andy S.H. To University
College London, Mullard Space Science Laboratory, Holmbury St. Mary, Dorking,
Surrey, RH5 6NT, UK

(Received April 3, 2023; Revised May 9, 2023; Accepted May 9, 2023)

###### Abstract

We report on the merging between the southern polar coronal hole and an
adjacent coronal dimming induced by a coronal mass ejection on 2022 March 18,
resulting in the merged region persisting for at least 72 hrs. We use remote
sensing data from multiple co-observing spacecraft to understand the physical
processes during this merging event. The evolution of the merger is examined
using Extreme-UltraViolet (EUV) images obtained from the Atmospheric Imaging
Assembly onboard the Solar Dynamic Observatory and Extreme Ultraviolet Imager
onboard the Solar Orbiter spacecraft. The plasma dynamics are quantified using
spectroscopic data obtained from the EUV Imaging Spectrometer onboard Hinode.
The photospheric magnetograms from the Helioseismic and Magnetic Imager are
used to derive magnetic field properties. To our knowledge, this work is the
first spectroscopical analysis of the merging of two open-field structures. We
find that the coronal hole and the coronal dimming become indistinguishable
after the merging. The upflow speeds inside the coronal dimming become more
similar to that of a coronal hole, with a mixture of plasma upflows and
downflows observable after the merging. The brightening of bright points and
the appearance of coronal jets inside the merged region further imply ongoing
reconnection processes. We propose that component reconnection between the
coronal hole and coronal dimming fields plays an important role during this
merging event, as the footpoint switching resulting from the reconnection
allows the coronal dimming to intrude onto the boundary of the southern polar
coronal hole.

††journal: ApJ††facilities: SDO/AIA, SDO/HMI, Solar Orbiter EUI, Hinode/EIS,
Hinode/XRT.††software: SunPy (SunPy Community et al., 2020), aiapy (Barnes et
al., 2020), SSW/IDL (Freeland & Handy, 1998), EISPAC
(https://github.com/USNavalResearchLaboratory/eispac)

##  1 Introduction

Coronal holes (CHs) are regions of relatively low density and temperature in
the solar corona, appearing dark when observed in Extreme-Ultraviolet (EUV)
and soft X-ray passbands (see review by Cranmer, 2009, and references
therein). Their main characteristic is the open magnetic field configuration,
which allows the plasma to escape the corona as the solar wind (Altschuler et
al., 1977; Levine et al., 1977). CHs generally have lifetimes of up to several
solar rotations (Harvey & Recely, 2002; Heinemann et al., 2018), and they can
be found in both polar regions (polar CH) and low-latitude regions (equatorial
CH). It is thought that polar CHs are formed gradually due to the result of
open flux accumulation and transport throughout the solar cycle, usually from
the poleward migration of open fields from lower latitude (Harvey & Recely,
2002; Lowder et al., 2017). In contrast, the formation of isolated, low-
latitude CHs was observed to be related to active region evolution during
either the emergence phase (Harvey & Hudson, 1998; Wang et al., 2010) or the
decaying phase (Karachik et al., 2010). CHs have been widely accepted as the
source of the fast solar wind, and CH boundaries have been proposed as one of
the sources of slow solar wind (see, e.g., Abbo et al., 2016; Cranmer et al.,
2017, and references therein). The fast solar wind streams are thought to
originate from open magnetic flux tubes concentrated at the supergranular
network boundaries, where they show the outflow velocities of $\sim$ 10 km s-1
as determined from spectroscopic observation using emission lines forming at
the base of the corona (Hassler et al., 1999; Tu et al., 2005). The slow solar
wind from CH boundaries, on the other hand, is thought to be associated with
the highly expanded open field near the CH edges (Wang & Sheeley, 1990), or be
driven by interchange reconnection at open-closed magnetic field boundaries
between CH and the surrounding regions (Antiochos et al., 2011; Abbo et al.,
2016; Aslanyan et al., 2022). Interchange reconnection also affects the
evolution of CH areas (Baker et al., 2007; Kong et al., 2018) and helps
maintain the quasi-rigid rotation of CHs (Wang & Sheeley, 2004; Yang et al.,
2011). Coronal jets, the small, transient collimated plasma ejections
frequently observed inside CHs, may also play an important role in supplying
mass and energy to the solar wind(Raouafi et al., 2023).

Some dark regions in the solar corona resemble coronal holes, but they have a
more transient nature and shorter lifespans. These features are historically
called transient coronal holes (Rust, 1983) but now are more commonly called
coronal dimmings (Hudson et al., 1996). They can be observed as regions of
reduced emission in soft X-ray and EUV passbands similar to CHs. Their
appearance is usually related to solar eruptions, such as coronal mass
ejections (CMEs) or solar flares (Bewsher et al., 2008; Dissauer et al.,
2018a, 2019). Coronal dimmings can be categorised into core and secondary
dimmings (Mandrini et al., 2007). Core dimmings have substantially reduced
emissions and often appear in pairs localised close to the eruption site. They
are interpreted as the footpoints of erupting flux ropes. Secondary dimmings,
in contrast, are more widespread and have a shorter lifespan and weaker
intensity reductions. They are understood to be the result of plasma escaping
along the expanding overlying field arcades above the flux ropes.

It is believed that the formation of coronal dimmings is due to density
depletion as the plasma escapes from the solar corona along the expanding CME
structure (Hudson et al., 1996; Thompson et al., 2000). This is supported by
the spectroscopic observation of strong plasma outflows of up to a hundred km
s-1 (e.g., Harrison & Lyons, 2000; Tian et al., 2012; Veronig et al., 2019).
Differential emission measure analysis (DEM) of coronal dimmings shows
significant plasma density decreases of up to 50-70 % (Cheng et al., 2012;
Vanninathan et al., 2018). It is also thought that coronal dimmings contain a
quasi-open magnetic field structure due to CMEs expansion, and the upflows may
then contribute to the solar wind (Harra et al., 2007; Attrill et al., 2010;
Lörinčík et al., 2021).

Interchange reconnection is also proposed to play an essential role in the
evolution of coronal dimmings (Attrill et al., 2006; Krista & Reinard, 2013)
and their recovery after eruption (Attrill et al., 2008). Liu et al. (2007)
and Gutiérrez et al. (2018) observed the formation of CHs from long-lived
coronal dimmings, suggesting that CHs can also form through an abrupt opening
of the magnetic field by solar eruptions (Heinemann et al., 2018).

In this paper, we report a rare event of the merging between a southern polar
CH and the coronal dimming resulting from a coronal mass ejection on 2022
March 18. This event occurred during the first perihelion in the science phase
of the ESA/NASA’s Solar Orbiter (SO; Müller et al., 2020) mission.
Concurrently, the Slow Solar Wind Connection Solar Orbiter Observing Plan
(Slow Wind SOOP; Yardley et al. 2023, submitted) was in operation during its
second remote sensing window (2022 March 17 – 22). This campaign used remote
sensing and in situ instruments onboard SO to investigate the origin of the
slow solar wind at open-closed magnetic field boundaries. The northwestern
edge of the south polar CH was chosen as the target during 2022 March 17–18,
which perfectly coincided with the time of the CME eruption and the merging
event.

The merging of coronal holes and coronal dimmings has only been reported a few
times (Nitta & Mulligan, 2017; Gutiérrez et al., 2018; Lörinčík et al., 2021;
Yardley et al., 2021b; Nitta et al., 2021), and has never been studied
spectroscopically due to their rare occurrence and transient nature.
Fortunately, this merging event was well captured by various spacecraft,
allowing us to analyse the merging in detail using high-resolution EUV and
X-ray observations alongside the coronal spectroscopic measurements. In this
paper, we investigate the changes in plasma dynamics inside the CH and the
coronal dimming due to the merging and propose a possible process behind the
merging of two open-field structures. The instrumentations and datasets used
in this study and the technique used to identify the boundaries of the CH and
coronal dimming are described in Section 2. The analysis results, including
the evolution of CH and coronal dimming, plasma dynamics across four different
timesteps during the merging event, and an example of a jet found inside the
merged region are presented in Section 3. Finally, we discuss the results and
their implications in Section 4 and summarise the manuscript in Section 5.

##  2 Observations

On 2022 March 18, a filament channel and the southern polar coronal hole were
observed on the solar disc. The filament partially erupted at $\sim$ 02:00 UT,
resulting in a coronal mass ejection and a pair of coronal dimmings. The
southeastern coronal dimming then expanded towards the southern polar CH due
to the expansion of CME. Around 08:15 UT, the coronal dimming reached the
northwestern tip of the CH boundary and appeared to merge with the CH, as
observed using EUV and X-ray passbands.

###  2.1 Coordinated Remote Sensing Observations

The merging and evolution of CH and coronal dimming were observed using EUV
images obtained from the Atmospheric Imaging Assembly (AIA; Lemen et al.,
2012) onboard Solar Dynamics Observatory (SDO; Pesnell et al., 2012), and the
Extreme Ultraviolet Imager (EUI; Rochus et al., 2020) onboard SO. SDO/AIA
continuously observes the full solar disk in seven EUV passbands, with a plate
scale of 0.6″ per pixel and 12 s cadence. EUI consists of two telescopes, the
Full Sun Imager (FSI) and High-Resolution Imager (HRI). FSI observes the full
solar disc with a resolution of 4.4″ per pixel and a cadence of 10 minutes.
HRI offers a smaller field-of-view (FOV) of 1000″$\times$1000″ but it has a
high spatial and temporal resolution with a plate scale of 0.492″ per pixel
and a cadence of 5 s. HRI was operating from 10:15 UT to 11:15 UT on 2022
March 18, capturing the merged region. Observations from AIA were downloaded
through the Virtual Solar Observatory (VSO) databases as level 1 data and
processed using the aiapy (Barnes et al., 2020) Python library. For the EUI
observations, we used calibrated level 2 EUI data from EUI data release 6
(Kraaikamp et al., 2023).

The line-of-sight (LOS) photospheric magnetograms, obtained from the
Helioseismic and Magnetic Imager (HMI; Scherrer et al., 2012) onboard SDO,
were used to study the magnetic field structure inside CH and coronal dimming.
The full disc magnetograms were obtained from the HMI front 4096 $\times$ 4096
pixel CCD camera, which measures the Zeeman splitting of the Fe I 6173.3 Å
absorption line. The temporal resolution is 45 s with a plate scale resolution
of 0.505″ per pixel. The photon noise level is approximately 7 G (Couvidat et
al., 2016). The HMI magnetograms were also downloaded from the VSO database
and were coaligned to match AIA observations.

In addition to the EUV images, we used coordinated observations from the
Extreme-Ultraviolet Imaging Spectrometer (EIS; Culhane et al., 2007) and the
X-Ray Telescope (XRT; Golub et al., 2007) onboard Hinode (Kosugi et al.,
2007). Both instruments were observing the boundary of the south polar CH on
2022 March 17–18, as planned in the Hinode Operation Plan 434. On the day of
the merging event, four EIS raster scans were obtained using two different EIS
studies. The first study was HPW021VEL240x512v2_b, a raster scan including a
large variety of spectral lines with a FOV of 240″$\times$ 512″, in 2″ slit
steps lasting 3 hr 34 min. This study was run at 01:22 UT and 11:52 UT.
Another study used was Atlas_60, a full CCD spectrum scan with a FOV of
120″$\times$160″. This study uses a 2″ slit with a total run time of 65 min.
This study was run at 07:13 UT and 09:50 UT. The Doppler and nonthermal
velocity can be derived from spectral data, with the velocity uncertainty of
$\sim$ 5 km s-1 for Doppler velocity and $\sim$ 20 km s-1 for nonthermal
velocity.

XRT was observing the southern polar CH boundary from 2022 March 17, 06:00 UT
to March 18, 19:40 UT and continuously provided soft X-ray observations using
the thin Al-poly filter, corresponding to a temperature of $\sim 10^{7}$ K.
The observations have a plate scale of 1″per pixel, and a FOV of
384″$\times$384″, except for the observation during 09:48–11:29 UT on March
18, where the resolution was resampled to 4″ per pixel, with a FOV of
2106″$\times$2106″.

EIS data files were downloaded from the UCL/MSSL database and were prepared
using standard EIS routines in the SolarSoftWare Library (SSW/IDL Freeland &
Handy, 1998). Spectra were corrected for instrumental effects, including slit
tilt, orbital variation, dark current, and warm/hot/dusty pixels. The XRT data
were obtained from Hinode/DARTS managed by ISAS/JAXA and were prepared using
the standard routine XRT_prep.pro, also available in SSW/IDL.

Figure 1: SDO/AIA 193 Å (left column) and EUI/FSI 174 Å (right column)
observations on 2022 March 18, capturing four stages of the merging between a
southern polar coronal hole (CH) and a coronal dimming. The CH and coronal
dimming boundaries are plotted in black and yellow, respectively. The FOVs of
the coordinated observations are shown in coloured rectangles. The white
arrows in panels a and b indicate the location of the filament channel, while
the flare loops are indicated by the white arrow in panel d. The intensity
maps have been enhanced using the Multiscale Gaussian Normalization method
(Morgan & Druckmüller, 2014). An animated version of this figure is available
as Fig1Movie_edit.mp4. The movie has a duration of 16 s and shows the
evolution of the polar CH and the coronal dimming as observed from SDO/AIA and
EUI/FSI from 00:00 UT to 14:00 UT.

###  2.2 Identifying Coronal Hole and Coronal Dimming Boundaries

Coronal holes and coronal dimmings are best observed at a coronal temperature
of $\sim 10^{6}$ K. Therefore, we used the 193 Å passband of SDO/AIA, and the
174 Å passband of EUI/FSI and EUI/HRI, to track the global structure of the CH
and coronal dimming. Note that SO was at an angle of 10∘ to the Sun-Earth
line, allowing us to observe this event from two different viewpoints.

The CH and coronal dimming boundaries were defined using SDO/AIA 193 Å
observations. For the CH, we implemented an intensity thresholding method,
with the threshold intensity chosen to be 35% of the solar disc median
intensity (Reiss et al., 2016; Hofmeister et al., 2017; Heinemann et al.,
2018). For the coronal dimming, on the other hand, we used base-ratio images
to enhance the relative changes in brightness (Dissauer et al., 2018b). The
SDO/AIA 193 Å observation at 01:30:04 UT, $\sim$30 minutes before the eruption
started, was chosen as the reference image. Subsequent images were divided by
the reference image to produce base-ratio images at each timestep. A
thresholding technique was then used to detect coronal dimmings, with a
threshold value of 0.35 (Dissauer et al., 2018b). After obtaining the initial
boundaries for the CH and coronal dimming, we applied a Gaussian filter to
each image to smooth the sharp edges and fill any holes inside the boundaries.
The resulting smooth-edge boundaries of CH and coronal dimming for each time
step were then projected onto other remote sensing observations.

##  3 Results

###  3.1 Evolution of Coronal Hole and Coronal Dimming

Figure 1 shows the evolution of the southern polar coronal hole and the
coronal dimming observed by SDO and SO at four different timesteps,
corresponding to the time at the middle of each EIS study111Note that the
distances of the individual spacecraft from the Sun meant that phenomena were
observed at different times by the different spacecraft. Hence, to avoid
confusion, we will use the time at Earth throughout this discussion.. The
boundaries of CH and coronal dimming are illustrated as black and yellow
contours, plotted over SDO/AIA 193 Å and EUI/FSI 174 Å observations. This
colour scheme is used throughout this paper. Each instrument’s FOVs are
displayed as coloured rectangles. The blue rectangle is the FOV of Hinode/EIS,
the red rectangle shows the FOV of Hinode/XRT, and the magenta rectangle shows
the FOV of EUI/HRI. An animated version of Figure 1 shows the evolution of the
polar CH and the coronal dimming from 00:00 UT to 14:00 UT.

From the SDO’s point of view, the northern tip of the southern polar CH was
close to the central meridian, with the northernmost edge at $y$ = -500″ . The
filament channel, indicated by the white arrows in Figure 1a-b, was located
$\sim$100″ west of the CH. It appeared to have a U-shaped structure spanning
at least half the size of the solar disk. Using photospheric LOS magnetograms
from HMI, we found that the CH had net negative polarity, the same polarity as
the eastern side of the filament channel.

Figure 2: Base-ratio image produced from dividing SDO/AIA 193 Å intensity at
03:05:04 UT by a reference image before the eruption (see text). Boundaries of
the coronal dimming and the CH are plotted in yellow and black, respectively.
The yellow arrow indicates the region between CH and coronal dimming where the
intensity enhancement (dark blue) is observed. Figure 3: SDO/AIA 193 Å
observation on 2022 March 21, showing the boundary of the southern polar
coronal hole three days after the merging occurred.

The filament partially erupted at the northern part of the U-shaped structure
at $\sim$02:00 UT, resulting in a CME, flare loops, and coronal dimmings. The
flare loops can be seen close to the central meridian from EUI/FSI 174 Å at
03:36 UT, as indicated by the white arrow in Figure 1d. The flare ribbons at
the footpoints of the flare loops can also be observed using 304 Å passband of
EUI/FSI (not shown here). Shortly after the eruption had started, the coronal
dimming started expanding in a southeastern direction towards the polar CH,
accompanied by the expansion of the flare loop footpoints. Figure 2 shows a
base ratio image which was used to define the boundaries of the coronal
dimming. The coronal dimming can be observed as an area of decreased
intensity, shown by red pixels. Figure 2 also shows the enhanced intensity
(dark blue pixels) at the region between the southern polar CH and a coronal
dimming, which are not evident in normal 193 Å intensity observations.

Figure 4: Plasma diagnostics of the coronal hole and coronal dimming at four
EIS time steps. From left to right, the columns show a) SDO/AIA 193 Å
intensity, b) XRT Thin Al-Poly filter intensity, c) Fe XII 192.39 Å Doppler
velocity map, d) Fe XII 192.39 Å nonthermal velocity map, and e) HMI line-of-
sight photospheric magnetogram. The boundaries of CH and coronal dimming are
shown in black and yellow, respectively. The colour scale of Doppler velocity
maps was set to [-10,10] km s-1. The nonthermal velocity maps are saturated at
[0,40] km s-1. HMI magnetograms are saturated at $\pm$ 50 G. The coloured
arrows indicate the interesting upflow regions discussed in Section 3.2.1.

The coronal dimming reached the northwestern edge of CH at $\sim$07:30 UT
before appearing to merge with the CH at $\sim$08:15 UT. Note that we define
the merging as the moment the CH and coronal dimming boundaries appear as a
single contour using the threshold method. Also, note that the stable part of
the filament (lower arc) was included in the merged region on the western side
beyond $x\sim$ 300″ due to its low intensity comparable to the CH in SDO/AIA
193 Å passband. However, the flare loop footpoints stopped their expansion at
around 11:30 UT, 3 hrs after the merging occurred. According to the standard
flare-CME model (CSHKP model, Carmichael, 1964; Sturrock, 1968; Hirayama,
1974; Kopp & Pneuman, 1976), the expansion of flare loop footpoints indicates
the ongoing reconnection process that simultaneously forms the flare loops and
builds more magnetic flux into the expanding CME, which results in the
expansion of coronal dimming. Therefore, we may infer that the coronal dimming
was still expanding after it merged with the CH until at least 11:30 UT.

In summary, the EUV and X-ray observations on 2022 March 18 show the evolution
of the coronal dimming and the southern polar CH as follows:

  * •

02:00 UT - The eruption started and the coronal dimming began its formation

  * •

03:36 UT - The flare loops can be seen from EUI/FSI 174 Å

  * •

08:15 UT - The coronal dimming merged with the southern polar CH

  * •

11:30 UT - The footpoints of flare loops stopped their expansion.

Figure 3 shows the merged region as seen from SDO/AIA 193 Å on 2022 March 21,
before the merged region moved away from the Earth-facing solar disc.
Interestingly, this figure shows that this merged region persisted for at
least three days after the merging occurred. The merged region then appeared
to shrink in size and later disappeared on 2022 March 22, leaving behind only
the southern part of the polar CH (the part of CH located approximately below
$y\sim$ -800″ ).

###  3.2 Spectroscopic Data Analysis

Hinode/EIS was pointing at the western boundary of the CH, where the merging
occurred. This allowed us to track and analyse the plasma dynamics at the time
before, during, and after the merging of the CH and the coronal dimming by
using four consecutive EIS raster scans. Although the merging of CHs and
coronal dimmings have been reported in previous literature, this is the first
time that such an event was observed spectroscopically for the entire process.

####  3.2.1 Overview For Each EIS Raster

Figure 4 shows the Doppler velocity and nonthermal velocity maps obtained from
fitting the Fe XII 192.39 Å spectral lines for four EIS rasters, along with
SDO/AIA 193 Å intensity, XRT thin Al-poly filter intensity and HMI line-of-
sight (LOS) magnetogram at the median time of each EIS scan (indicated on the
right axis of each row). We will use the median time to represent the
observation time of each study throughout the paper. The Fe XII 192.39 Å
spectral line was used because it is unblended, and the reference wavelength
is close to the peak emission and temperature response of the SDO/AIA 193 Å
observation. All observations were aligned to the SDO/AIA 193 Å intensity
maps, with the CH and coronal dimming boundaries then overplotted.

The first raster was taken at $\sim$03:05 UT, associated with the expansion of
coronal dimming toward the east as the filament eruption was in progress. In
this raster, we can see a large region of strong blueshift in the Doppler
velocity map (see panel 1c in Figure 4), indicating a strong plasma upflow
inside the coronal dimming region. This structured strong plasma upflow inside
the coronal dimming has been reported in previous literature and usually has
been interpreted as a result of rapid plasma evacuation onto expanding CME
structure (see, e.g., Harra et al., 2007; Attrill et al., 2010; Tian et al.,
2012; Veronig et al., 2019).

During the second raster at $\sim$07:45 UT, EIS observed the southeastern edge
of the coronal dimming. During that time, the coronal dimming was located
close to the southern polar CH boundary. Several bright points, previously
located in the quiet Sun region, were found inside the boundary of coronal
dimming (see panel 2a, 2b in Figure 4). There are two significant upflows
regions in the Doppler velocity map, indicated by pink and orange arrows in
panel 2c of Figure 4. The first region (pink arrow) is located between two
bright points at (260″, -680″). Interestingly, this region had no observable
signature in SDO/AIA 193 Å and XRT observation, despite clear blueshift and
enhanced nonthermal velocity. The second region (orange arrow) is located at
(170″, -630″), close to a bright point and the interface between the CH and
coronal dimming boundaries. Similarly, the source for this upflow is also
uncertain.

The third EIS raster observed the newly merged at $\sim$10:20 UT. A clear
signature of a thin collimated jet can be seen at (190″, -650″), indicated by
two pink arrows in panels 3a and 3c of Figure 4. This feature can be observed
as a thin corridor of blueshifted plasma in the EIS Doppler velocity map and a
collimated thread-like funnel in SDO/AIA 193 Å observation. The base of upflow
is located next to a bright point, which is the same bright point located next
to the observed upflow indicated by the orange arrow in panel 2c of Figure 4.

Figure 5: EIS Fe XII 192.39 Å Doppler velocity maps obtained from EIS studies
centred at $\sim$03:05 UT and $\sim$13:35 UT. The orange boxes indicate the
approximate locations of regions where we averaged the spectral data to obtain
the plasma dynamics shown in Table 1.

The fourth raster was run at $\sim$13:35 UT, 5 hours after the merging
occurred. The Doppler and nonthermal velocity maps show that plasma flow
became more disordered at this time, as indicated by the enhanced nonthermal
velocity and the mixture of plasma upflow and downflow inside the merged
region. However, this may be partially due to instrumental effects such as
warm pixels. A relatively structured upflow region was observed at (250″,
-650″), indicated by a pink arrow in panels 4a and 4c in Figure 4. It seems to
be related to a breakout jet resulting from a mini-filament eruption at the
bright point that occurred at $\sim$13:30 UT.

####  3.2.2 Spatially Averaged Pixel Analysis

We use the spatially averaging method described in EIS Software Note No.17
222Access through
https://vsolar.mssl.ucl.ac.uk/JSPWiki/Wiki.jsp?page=EISAnalysisGuide to obtain
important plasma dynamic parameters across the four EIS rasters in several
regions of interest. In doing this, we sacrifice the spatial resolution to
enhance the signal-to-noise ratio of the low-emission coronal hole and coronal
dimming.

First, we select four regions of interest, denoted as Box 1–4 in Figure 5. Box
1 is the thin upflow corridor at the north tip of the CH that transcended the
boundary into the quiet Sun region. Box 2 is the strong upflow region inside
the coronal dimming, Box 3 is the quiet Sun region between the CH and the
filament channel before the merging event. Lastly, Box 4 is the jet’s location
observed in the EIS raster at 10:20 UT. The approximate locations of these
boxes are shown in Figure 5. Next, we treat each box as a single pixel by
averaging spectral data inside the box. We then derive the plasma parameters
by fitting a Gaussian function to the averaged spectral data. Note that in
some cases, we choose smaller regions at the same location with the boxes
shown in Figure 5 to obtain a better spectral fit.

Box Number (Feature) | EIS raster | Doppler Velocity (km s-1) | Nonthermal Velocity (km s-1) |  | Secondary
---
Component
Box 1 (CH Edge) | 03:05 UT | -18 | 28 | No
13:35 UT | -22 | 31 | No
Box 2 (Coronal Dimming) | 03:05 UT | -26 | 20 | Yes
13:35 UT | -16 | 33 | No
Box 3 (Quiet Sun) | 03:05 UT | 0 | 25 | No
13:35 UT | -3 | 31 | No
Box 4 (Jet) | 03:35 UT | -9 | 27 | No
07:45 UT | -6 | 28 | No
10:20 UT | -30 | 24 | Yes
13:35 UT | -23 | 34 | No

Table 1: Results from EIS analysis using the averaged spectral data inside
four boxes of interest. The table shows the Doppler and nonthermal velocities
for each box at different times and whether the spectra have significant
signatures of the secondary components in the line profile. Note that the
merging occurred at $\sim$08:15 UT. Figure 6: Fe XII 192.39 Å emission line
profiles obtained from the spatially averaged analysis for coronal dimming
upflow (Box 2 in EIS raster at 03:05 UT; left panel) and jet upflow (Box 3 in
EIS raster at 10:20 UT; right panel). Both line profiles show the signatures
of blue-wing enhancement as indicated by the double gaussian fits, blue lines
for the primary component and red lines for the secondary component. The
calculated LOS velocity and nonthermal velocity of the secondary component are
displayed in each panel.

Table 1 shows the results obtained from this spatially averaged pixel analysis
using data from the Fe XII 192.39 Å spectral line on all available four EIS
rasters. Boxes 1, 2 and 3 are located outside the FOVs of the EIS rasters at
07:45 UT and 10:20 UT. Therefore, only Box 4 is observed on every raster.

In Box 1 (CH edge), we observed the upflow velocity of $\sim$20 km s-1 and the
nonthermal velocity of $\sim$30 km s-1. These values are similar to the
typical value of plasma dynamics inside CH observed using coronal spectral
line (e.g., Harra et al., 2015; Fazakerley et al., 2016).

Boxes 3 and 4 are associated with the quiet Sun region before the merging,
which later became the merged region of CH and coronal dimming and showed
significant changes in plasma dynamics. Box 3 shows an increase in LOS
velocity, as evident from the stronger redshift and blueshift pixels inside
the box. The upflow profile changes from no upflow (on average) to upflow of
$\sim$3 km s-1. It also shows an increase in nonthermal velocity. The increase
may be due to a change in the magnetic environment inside the region; the
initially closed field region became more open because of the merging between
CH and coronal dimming. The change in LOS velocity inside Box 4 is similar to
Box 3, showing a more significant upflow when comparing between EIS raster at
03:05 UT and 13:35 UT, with a similar increase in nonthermal velocity.
Furthermore, plasma dynamics remained almost the same in the EIS raster at
03:05 UT and 07:45 UT before changing drastically at the EIS raster at 10:20
UT, when there was a thin, collimated jet outflow inside the box.

Lastly, only Box 2 shows a decrease in LOS upflow velocity after the merging.
The initial high upflow velocity in coronal dimming may be due to the plasma
evacuation as the CME expanded. After the merging, the upflow velocity
decreased to a level seen in the CH, while the nonthermal velocity increased
similarly to other boxes.

By examining each line profile in detail, we see that some profiles have
noticeable enhancements in the blue-wing region. This enhancement suggests
that there are at least two components of unresolved plasma upflow along the
LOS (Tian et al., 2021); namely, the background plasma flow (primary
component) and high-speed upflow (secondary component). The strong secondary
components have previously been observed in the upflows from active region
boundaries (e.g., Baker et al., 2021; Yardley et al., 2021a) and coronal
dimming (Tian et al., 2012; Veronig et al., 2019). They can be identified by
fitting a double gaussian instead of a single gaussian function to the
profile. In our analysis, the averaged pixels from Box 2 in EIS raster at
03:05 UT and Box 4 in EIS raster at 10:20 UT show the significant secondary
component in line profiles, i.e., the secondary component intensity is more
than 10% of the primary component (Tian et al., 2021). These line profiles
correspond to the coronal dimming (Box 2) and the jet inside the merged region
(Box 4). Both line profiles are shown in Figure 6; the coronal dimming in the
left panel and the jet in the right panel. We also see comparable blue-wing
enhancement in both regions in Fe XII 195.12 Å and Fe XIII 202.04 Å line
profiles.

Figure 7: A coronal jet inside the newly merged region as observed from
SDO/AIA 193 Å passband (left), HMI magnetogram (centre), and EUI/HRI 174 Å
passband (right) at 10:35 UT. The coloured arrows show the locations of the
jet, the bright point and the adjacent small coronal loops. The contour of jet
upflow observed in the EIS study centred at 10:20 UT is overplotted in dark
blue. The green and yellow contours correspond to the magnetic field strengths
of -50 and 50 G, respectively. An animated version of this figure is available
as Fig7Movie_edit.mp4. The movie has a duration of 24 s and shows the
evolution of the jet as observed from SDO/AIA, SDO/HMI and EUI/HRI from 10:15
UT to 11:15 UT.

For the case of coronal dimming, the line profile looks like a single gaussian
function with a bump on its blue-wing tail, indicating that the secondary
upflow was weak but had a very high velocity. The secondary component has a
LOS velocity of -187 km s-1, higher than the reported value of CME upflows in
the range of [50–150] km s-1 (Tian et al., 2012; Veronig et al., 2019). We
interpret that the high-speed plasma upflow is a result of impulsive plasma
outflow along the expanding magnetic field structure (Harrison & Lyons, 2000;
Harra et al., 2007; Veronig et al., 2019) rather than the plasma refilling the
coronal dimming after the eruption (Tian et al., 2012) due to the extended
coronal dimming lifetime. The nonthermal velocity is 19 km s-1, lower than the
primary upflow. For the jet inside the merged region, the secondary component
has a LOS velocity of -64 km s-1. It also has a nonthermal velocity of 32 km
s-1, higher than the primary upflow. The secondary components for both
features can be observed in only one EIS raster, indicating the transient
nature of high-speed plasma outflows.

###  3.3 A Coronal Jet at 10:30 UT

X-ray and EUV observations show several coronal jets and brightenings of
bright points during this merging event, especially in the interface region
between the coronal hole and coronal dimming (see panels 2a, 2b, 3a and 3b of
Figure 4). These bright points initially resided in the quiet Sun region
between the CH and the filament channel, which later turned into the merged
region.

An example of a coronal jet is shown in Figure 7. We selected this jet as an
example because it was well observed in SDO/AIA, EUI/HRI and Hinode/EIS
Doppler velocity map. Figure 7 displays a thin-collimated beam of plasma
observed by SDO and SO at $\sim$10:35 UT. The jet is located around (180″,
-650″) in helioprojective coordinates as seen from SDO (blue arrows), close to
the boundary of the post-merging CH and the bright point (red arrows). It also
corresponds to the thin upflows corridor observed in the EIS raster at 10:20
UT (indicated by a pink arrow in panel 3c of Figure 4), which is shown as the
blue contours in Figure 7. The jet became visible in SDO/AIA 193 Å at
$\sim$10:00 UT. Note that the jet formation might start earlier but is not
noticeable due to the brighter background emission. The jet initially had a
weak emission and consisted of several thin upflows threads, with the jet’s
base adjacent to the eastern edge of the bright point. It then appeared
brighter and became more collimated into a thin beam around 10:25 UT. Besides
the coronal jet, we also observed brightenings of coronal loops in the bright
points and the adjacent loops under it (pink arrows in Figure 7) at
$\sim$10:30 UT. Finally, the jet slowly dispersed and completely disappeared
at $\sim$11:00 UT, marking a visible lifetime of $\sim$1 hr.

The jet can be classified as a standard jet due to its similar configuration
(long-thin spire) to the model proposed by Shibata et al. (1992) and the lack
of emission in the 304 Å passband (cf. Long et al., 2023). Using the
spectroscopic analysis discussed earlier, this jet had a LOS velocity of -64
km s-1 with a nonthermal velocity of 32 km s-1 (see right panel on Figure 6).

This jet is not clearly visible against brighter background emission in
EUI/HRI 174 Å observations. However, we notice that the background emissions
seem to have an outflowing motion from the bright point that the jet
originated from, suggesting that they were coronal plumes, hazy ray-like
transient structures commonly found in CH. Coronal plumes can be frequently
observed using white light and SDO/AIA 171 Å passbands and may also contribute
to supplying mass and energy to solar wind (see review by Poletto, 2015, and
references therein). The high resolution of EUI/HRI can resolve the motion of
closed loops inside the bright point and surrounding area, which can be
related to the reconnection process powering the jet. The observations also
reveal numerous small-scale, short-lived brightenings that are similar to
‘campfires’ (Berghmans et al., 2021) observed in the quiet Sun region.

The photospheric magnetogram (the middle panel in Figure 7) shows the minority
polarity (positive, yellow contour) of the bright point surrounded by
background majority polarity (negative, green contour) in the merged region.
In our scenario, this configuration resembles dipoles embedded in a uniform
open field, which is often associated with a 3D fan-spine structure of closed
loops with a domed separatrix surface that separates closed and open magnetic
field regions (e.g., Kumar et al., 2019). It is also frequently used as the
initial setup for simulating coronal jet formation (Yokoyama & Shibata, 1995;
Pariat et al., 2009; Wyper et al., 2018). In these setups, the interchange
reconnections between the open background field and closed loops are
responsible for allowing the reconnection exhaust to escape as jets along the
newly open field lines, which is compatible to our observations.

Note that there were multiple brightenings of bright points and jets inside
the merged region other than this jet in Figure 7. These jets can be different
from the example jet discussed above. For example, the jet observed at
$\sim$13:30 UT (indicated by pink arrows in panel 4a and 4c in Figure 4) was
likely to be a breakout jet corresponding to a mini-filament eruption from the
bright point, as the EUV observations show the eruption of a small coronal
loop with the brightenings observable in 304 Å passband.

##  4 Discussion

Our observations show that the coronal dimming intruded the boundary of the
southern polar coronal hole, resulting in the merging of two open field
structures. The merged region persisted for over 72 hr (see Figure 3),
significantly longer than the typical coronal dimming lifetime of 3–12 hr
(Reinard & Biesecker, 2008). The merged region then shrank and disappeared
alongside the northern part of CH, suggesting that the coronal dimming became
indistinguishable from the polar CH. If there were no interactions between the
coronal dimming and the CH, the northern part of the CH should still exist
following the disappearance of the coronal dimming. The spectroscopic
observations from Hinode/EIS show that the upflow speed inside the coronal
dimming region decreased to a roughly similar level with the CH after the
merging, with no significant sign of a secondary Gaussian component,
indicating an absence of superposed high-speed plasma upflow associated with
the coronal dimming (see Table 1). The nonthermal velocities in every region
included in the analysis were enhanced after the merging, further suggesting
that there were underlying mechanisms contributing to the broadening of the
line profile.

###  4.1 Component Reconnection as Merging Mechanisms

Figure 8: A 2D diagram illustrating the merging process of a coronal hole (CH,
black) and a coronal dimming (CD, orange). (a) Initially, the CH and the
coronal dimming were further apart. Coronal dimming field lines (orange lines)
were at an angle to the CH field lines (black lines). (b) The coronal dimming
then expanded towards the CH boundary. The field lines from both structures
reconnected at the interface region. The zoomed-in view of the reconnection
site is shown in the red square. The CH and coronal dimming magnetic field
lines had an anti-parallel component, allowing component reconnection to
occur. (c) The footpoints of CH and coronal dimming field lines were switched
after the reconnection, as indicated by newly reconnected field lines (blue
and pink). The mixture of plasma upflows and downflows observed after the
merging may also be due to plasma outflows from the reconnection site.

We propose that the main interaction involved in this merging process is the
footpoint switching between open field lines in the CH and coronal dimming.
Since both features had dominant negative polarity, we suggest that component
reconnection was responsible for this process. Component reconnection is
magnetic reconnection in which the magnetic field lines are not anti-parallel
to each other, but there is still some angle between reconnecting field lines
(i.e., one component of the field lines is anti-parallel). This type of
reconnection can occur in unipolar magnetic field regions and can be found in
various places and scales throughout the solar atmosphere (e.g., Parker, 1983;
van Driel-Gesztelyi et al., 2012; Rappazzo et al., 2012; Tenerani et al.,
2016; Antolin et al., 2021; Chen et al., 2021).

In this scenario, the main driver for the component reconnection between the
CH and the coronal dimming is the extensive lateral expansion of the CME.
Figure 2 shows the brightening at the region between coronal dimming and the
southern polar CH in the base-ratio image. We speculate that this brightening
may be a result of the plasma compression inside the overlying field arcade
above the expanding CME since the overlying arcade may not be able to expand
further east due to the open field of the coronal hole. Hence, the brightening
implies that the lateral expansion of CME is powerful enough to compress the
coronal plasma and potentially drive the reconnection process.

Figure 8 shows a diagram explaining the component reconnection during our
merging event. Initially (panel a), the CME resulted in a coronal dimming
which contained the newly open field. The coronal dimming field lines (orange
lines) were at an angle with the field lines of nearby polar CH (black lines)
due to the expansion of the CME. The boundaries of coronal dimming then
reached the CH (panel b). Since at least one component of magnetic field lines
from each region was anti-parallel, the component reconnection between coronal
dimming and coronal hole magnetic fields could take place at the boundary
interface, as shown in the diagram inside the red rectangle in panel b.

Reconnection reconfigured the field lines and resulted in footpoint switching
(panel c); the post-reconnection coronal dimming field (pink line) is rooted
in the CH, and the CH field (blue line) is rooted in the coronal dimming. This
made the boundaries of both regions overlap with each other, resulting in the
merging and allowing the coronal dimming to intrude on the boundary of the CH.

This scenario can explain why the coronal dimming was still able to expand
after the merging occurred (as indicated from the expansion of flare loops
footpoints in Figure 1) since the ongoing footpoint switching allows the
coronal dimming to keep expanding further inside the coronal hole. The
footpoint switching might also distribute the CH fields throughout the merged
region, which may help make the CH and coronal dimming indistinguishable and
maintain the merged region’s structure. The plasma outflow due to the
component reconnection is a plausible explanation for the mixture of
disordered plasma upflow and downflow and the enhancement of nonthermal
velocities observed after the merging in EIS Raster 13:35 UT (see panel 4c and
4d of Figure 4).

We hypothesise that the component reconnection started at a higher altitude
and moved to a lower height as the merging progressed. The reconnection
occurred continuously until the expansion of coronal dimming stopped, at which
point there was no further movement of field lines that could drive the
reconnection process.

###  4.2 Jets and Upflows Inside Coronal Dimming and the Merged Region

Using the broadband EUV and spectroscopic observations, we found numerous
signatures of coronal plasma outflows throughout the merging processes, either
as coronal jets in EUV observations or blueshift (upflow) regions in the
Doppler velocity maps.

These events were more apparent after the merging, which could be due to the
darker background. However, we believe that the main reason is the enhanced
reconnection due to the change of the magnetic configuration caused by the
expanding coronal dimming and the subsequent merging with CH. Before the
eruption, these bright points resided in the quiet Sun region between the
southern polar CH and the filament channel, where the magnetic configuration
mainly consisted of closed magnetic loops. The erupting CME opened up the
field lines, resulting in the observed coronal dimming, which later expanded
to the region where bright points were located. This drastically changed the
magnetic configuration from a closed environment to an open one, allowing more
interchange reconnections and jet outflows.

In our observations, the open fields from the coronal dimming expanded across
the bright point where the jet occurred. This scenario is analogous to the
simulation setup for studying jets and bright point brightenings due to the
moving magnetic elements in a CH by Wyper et al. (2018). In the first
configuration in the study, they simulated a moving embedded minor polarity
inside the background field of the net majority polarity, inducing the shear
flow at the separatrix surface, which is similar to our scenario in the moving
minority polarity frame. The results from the simulation indicated quasi-
periodic weak bursts of interchange reconnection, which resulted in the
brightening of closed loops and weak outflow. The outflow resembles a standard
jet (i.e., long thin spire with no emission in the 304 Å passband) with an
outflow speed of around 60 km s-1.

An example jet in Figure 7 has a LOS velocity of -64 km s-1 and has no
signature in 304 Å passband, which agrees with the results of the simulation
by Wyper et al. (2018). Therefore, we believe that this jet corresponds with
the simulation setup. However, the main differences between our events and the
simulation are the driving force of the magnetic field movement and the angle
of the background open field. The simulation drove the magnetic elements
through the vertical field using photospheric flows, in contrast with the
angled field passing through bright points due to the expansion of coronal
dimming seen here.

Although we reported the source of several upflows in the results, there are
still some upflows whose sources were less certain, such as the two upflow
regions inside the coronal dimming region observed in the EIS raster 07:45 UT
(see the coloured arrows in panel 2c of Figure 4).

The first upflow region (pink arrow) is located well inside the coronal
dimming, showing a strong blue-shifted plasma with no EUV and X-Ray emission
counterpart. One possible explanation is that these outflows are escaping
plasma alongside expanding structure caused by the eruption, similar to CME-
induced upflows observed in the EIS raster at 03:05 UT (panel 1c of Figure 4).
However, due to its localised shape and position close to two bright points,
the upflow could also be a structure resembling a ‘dark jet’ inside CHs
discovered in a study by Young (2015). The second outflow region (orange
arrow) is located at the interface region between the CH and coronal dimming
and close to the bright point, which was the source of the jet in Figure 7.
Hence, the upflows could also be a jet originating from the bright point.
However, the brighter quiet Sun background in both EUV and X-ray observation
made jet identification challenging. This upflow can also be caused by
reconnection at the boundary region between CH, coronal dimming, and quiet
Sun, either through interchange reconnection or component reconnection between
open fields.

###  4.3 Relationship between Coronal Holes and Coronal Dimmings

Coronal dimmings are often called transient coronal holes due to their similar
appearance. However, they can be distinguished by their difference in lifetime
and formation process. Coronal dimmings arise from solar eruptions and have a
lifetime on a scale of hours, while CHs are thought to form more gradually and
usually persist for several Carrington rotations. However, recent observations
indicate that solar eruptions could be another way to form CHs (Liu et al.,
2007; Gutiérrez et al., 2018; Heinemann et al., 2018), suggesting that some
CHs may actually be coronal dimmings with unusually long lifetimes. Following
this train of thought, we may infer that long-lived coronal dimmings and CHs
are almost identical. Our observations show that the coronal dimming (merged
region) persisted for more than 72 hr after its formation and became
indistinguishable from the CH after the merging. This provides another example
of a CH arising from coronal dimming and supports the idea that coronal
dimmings and CHs are similar structures.

The recovery of coronal dimmings is thought to be facilitated by interchange
reconnection with the surrounding quiet Sun regions (Attrill et al., 2008),
which disperses the concentration of open field inside the coronal dimming and
allows plasma to fill in the dimming region. Therefore, the speed of coronal
dimming recovery may be related to the surrounding magnetic field environment.
In a unipolar region, the coronal dimming recovery is expected to be slower
due to fewer closed loops to reconnect with open fields. In our observation,
the coronal dimming was located near the southern polar CH and later merged
with each other. Since the coronal dimming and the CH had the same net
magnetic polarity, this created a large, strong unipolar region which could
slow the dispersion of the open fields to surrounding quiet Sun regions,
resulting in an unusually long lifetime of coronal dimming that later became
indistinguishable from CH. Several mergers of CHs and coronal dimmings and the
conversions from coronal dimming to CH were reported by Gutiérrez et al.
(2018), and all of them were found in unipolar magnetic environments.
Therefore, we speculate that if eruptions occur near unipolar regions, the
resulting coronal dimmings will likely have longer lifetimes and could later
turn into isolated CHs or merge with a neighbouring CH.

Like CHs, coronal dimmings have been suggested as a potential source of the
solar wind due to observed outflows and open-field configuration (Thompson et
al., 2000; Harra et al., 2007; McIntosh et al., 2007). Lörinčík et al. (2021)
reported evidence of solar wind from a coronal dimming as continuous outflow
funnels were observed in 171 Å and 193 Å passbands of SDO/AIA. They lasted
more than five hours and had a velocity of up to 70–140 km s-1. The derived
outflow speed agrees with the coronal dimming upflow speed obtained from
spectroscopic observation in our work and previous literature (Tian et al.,
2012; Veronig et al., 2019), and has a similar outflow pattern and velocity
with neighbouring CH. Our result suggests that coronal jets and coronal plumes
can originate from bright points inside the coronal dimming and could
contribute to solar wind. The component reconnection between open fields,
which could occur during our merging scenario, can also be an important
process for solar wind acceleration (Viall & Borovsky, 2020). However, post-
observation connectivity analysis using the magnetic connectivity tool
(Rouillard et al., 2020) reveals that it was unlikely that SO was magnetically
connected with the western boundary of CH during the merging event. Therefore,
the solar wind streams arising from this region were likely not observed by
the spacecraft.

##  5 Conclusion

In this paper, we report and analyse the merging of the southern polar coronal
hole with a coronal dimming resulting from a coronal mass ejection on 2022
March 18 during the Slow Wind SOOP coordinated observation campaign. The
coordinated observation allowed us to analyse the merging process of two open-
field structures in detail by combining high spatial and temporal resolution
EUV and X-ray observations with the coronal spectroscopic data. The main
results and discussion can be summarised as follows.

  1. 1.

The merging of the polar coronal hole and a coronal dimming resulted in the
merged region that persisted for more than 72 hours, significantly longer than
a typical lifetime of coronal dimmings. The cotemporal disappearance of the
northern part of the coronal hole alongside the coronal dimming indicates that
the coronal dimming and the coronal hole became indistinguishable from one
another.

  2. 2.

The plasma dynamics were observed to become more disordered after the merging
of the coronal hole and the coronal dimming, as evident in the mixing of
upflow and downflow throughout the region and the overall enhancement of
nonthermal velocities. Using a spatially averaged pixel analysis, we find that
the upflow profile inside the coronal dimming and quiet Sun regions became
more similar to the coronal hole upflow after the merging. In particular, the
upflow velocity inside the coronal dimming changed from the order of hundred
km s-1 initially to the order of ten km s-1 after the merging occurred.

  3. 3.

Component reconnection is considered the primary driving process in the
merging of these two structures. The open fields from the coronal hole and the
coronal dimming reconnect and exchange their footpoints, allowing the coronal
dimming to intrude into the coronal hole region and making the coronal hole
and the coronal dimming indistinguishable. The mixture of plasma upflow and
downflow observed after the merging may also be due to the reconnection.

  4. 4.

Several bright point brightenings and coronal jets are found inside the
coronal dimming and the merged region throughout the observation period,
suggesting ongoing reconnection processes. A study of the coronal jet inside
the merged region shows that this jet may be driven by the interchange
reconnection between the closed loops in a bright point and the expanding open
fields in the coronal dimming.

Coronal dimmings are traditionally called transient coronal holes due to their
resemblance to one another. However, we provide evidence that they are also
closely related, as coronal dimmings can merge with coronal holes and become
the same structure. Their similarities also suggest that coronal dimming may
play an important role in the formation of coronal holes. Further studies are
needed to understand what happens during the merging of two open-field
structures and the potential consequences in the heliosphere, such as precise
magnetic field extrapolation during the merging, the study of plasma
composition evolution throughout the merging process, and the in-situ
observation of the outflows from the merged region. Our study also emphasises
the importance of coordinated solar observations as a powerful means to reveal
the mysteries of the Sun. By combining several observations from various
spacecraft, we can improve our understanding of the coronal features, their
magnetic field structures and their influence towards the heliosphere in the
form of CME or solar wind.

We are grateful to Lidia van Driel-Gesztelyi, Gherado Valori and David Brooks
for their helpful discussions which improves the content of this paper. Solar
Orbiter is a space mission of international collaboration between ESA and
NASA, operated by ESA. The EUI instrument was built by CSL, IAS, MPS,
MSSL/UCL, PMOD/WRC, ROB, LCF/IO with funding from the Belgian Federal Science
Policy Office (BELSPO/PRODEX PEA 4000134088); the Centre National d’Etudes
Spatiales (CNES); the UK Space Agency (UKSA); the Bundesministerium für
Wirtschaft und Energie (BMWi) through the Deutsches Zentrum für Luft- und
Raumfahrt (DLR); and the Swiss Space Office (SSO). SDO data are courtesy of
NASA/SDO and the AIA, EVE, and HMI science teams. Hinode is a Japanese mission
developed and launched by ISAS/JAXA, collaborating with NAOJ as a domestic
partner, NASA and STFC (UK) as international partners. Scientific operation of
the Hinode mission is conducted by the Hinode science team organized at
ISAS/JAXA. This team mainly consists of scientists from institutes in the
partner countries. Support for the post-launch operation is provided by JAXA
and NAOJ (Japan), STFC (UK), NASA (USA), ESA, and NSC (Norway). N.N. is
supported by STFC PhD studentship grant ST/W507891/1 and UCL Studentship.
D.M.L. is grateful to the Science Technology and Facilities Council for the
award of an Ernest Rutherford Fellowship (ST/R003246/1). D.B. is funded under
STFC consolidated grant number ST/S000240/1. S.L.Y. would like to thank NERC
for funding via the SWIMMR Aviation Risk Modelling (SWARM) Project, grant
number NE/V002899/1, and STFC via the consolidated grant number STFC
ST/V000497/1. A.W.J. was supported by a European Space Agency (ESA) Research
Fellowship and acknowledges funding from the STFC consolidated grant
ST/W001004/1.

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