# Observations and detectability of young Suns’ flaring and CME activity in
optical spectra

M. Leitzinger,1 P. Odert,1 R. Greimel2
1Institute of Physics, Department for Astrophysics and Geophysics, University
of Graz, Universitätsplatz 5, 8010 Graz, Austria
1RG Science, Schanzelgasse 17, 8010 Graz, Austria E-mail:
<EMAIL_ADDRESS>
(Accepted XXX. Received YYY; in original form ZZZ)

###### Abstract

The Sun’s history is still a subject of interest to modern astrophysics.
Observationally constrained CME rates of young solar analogues are still
lacking, as those require dedicated monitoring. We present medium resolution
optical spectroscopic monitoring of a small sample of bright and prominent
solar analogues over a period of three years using the 0.5m telescope at
observatory Lustbühel Graz (OLG) of the University of Graz, Austria. The aim
is the detection of flares and CMEs from those spectra. In more than 1700
hours of spectroscopic monitoring we found signatures of four flares and one
filament eruption on EK Dra which has been reported in previous literature,
but we complementarily extended the data to cover the latter phase. The other
stars did not reveal detectable signatures of activity. For these non-
detections we derive upper limits of occurrence rates of very massive CMEs,
which are detectable with our observational setup, ranging from 0.1 to 2.2
day-1, but these may be even smaller than the given rates considering
observational biases. Furthermore, we investigate the detectability of
flares/CMEs in OLG spectra by utilizing solar 2D H$\alpha$ spectra from Mees
solar observatory. We find that solar-sized events are not detectable within
our observations. By scaling up the size of the solar event, we show that with
a fractional active region area of 18% in residual spectra and 24% in
equivalent width time series derived from the same residuals that solar events
are detectable if they had hypothetically occurred on HN Peg.

###### keywords:

stars: activity – stars: chromospheres – stars: flare – stars: late-type –
stars: individual: EK Dra – Sun: Coronal Mass Ejections

††pubyear: 2015††pagerange: Observations and detectability of young Suns’
flaring and CME activity in optical spectra–6

##  1 Introduction

The Sun has a 4.6 Gyr long history which was subject to numerous
investigations. The “Sun in time” program (see e.g. Dorren & Guinan, 1994;
Güdel, 2007) was founded to investigate the Sun’s history in great detail. The
radiation environment was reconstructed from X-rays (e.g. Dorren et al., 1995;
Güdel et al., 1997b; Telleschi et al., 2005; Guinan, 2017), EUV (e.g. Güdel et
al., 1997a; Tu et al., 2015), FUV (e.g. Guinan et al., 2003), UV (e.g. Dorren
& Guinan, 1994; Dalton et al., 2019), optical (e.g. Messina & Guinan, 2002) to
radio (e.g. Güdel et al., 1994; Villadsen et al., 2014; Fichtinger et al.,
2017). The spectral energy distributions of the Sun in time has been inferred
(Ribas et al., 2005; Claire et al., 2012) and also the solar wind in time has
been investigated (Ó Fionnagáin & Vidotto, 2018).
Every study focusing on solar analogues of different age may be attributed to
the idea of the “Sun in time” program. Transient activity phenomena like
flares and CMEs of the young Sun can be characterized with a significant
observational effort only as those are detectable via time series observations
which require much observing time. Here, especially the CME environment of the
young Sun remains still relatively unknown. However, flare frequency
distributions, as well as flare power laws depending on the stars X-ray
luminosity of young solar analogue stars and others have been presented by
Audard et al. (2000). Based on these power laws Odert et al. (2017) have
deduced relations to estimate stellar CME occurrence rates. Prior to Odert et
al. (2017), Aarnio et al. (2012) established a methodology to relate solar
flare/CME relations with stellar flaring relations to infer stellar CMEs and
their parameters. Drake et al. (2013) applied a similar approach and
identified the problem of the unknown stellar flare-CME association rate, as
extrapolating to higher energies while using solar relations leads to
unrealistic high energy requirements which have been not observed yet. Osten &
Wolk (2015) assumed energy partition between bolometric flare radiation and
kinetic energy of the associated CME. These authors found mass loss rates
comparable to previous studies. To explain energy budget problem discussed in
Drake et al. (2013), Odert et al. (2017) suggested then that probably the
whole flare-CME association rate may shift to larger energies.
Stellar flares are a subject of ongoing research going back to the first half
of the last century where stellar flares have been detected using ground-based
observations (Joy & Humason, 1949; Luyten, 1949, followed by numerous
studies). With satellite missions such as the Microvariability and
Oscillations of Stars Telescope (MOST), Convection, Rotation and planetary
Transits (CoRoT), Kepler and now with the Transiting Exoplanet Survey
Satellite (TESS) and in the near future also with the PLAnetary Transits and
Oscillation of stars mission (PLATO) long-term photometric measurements were
and will be accessible. This enabled statistical investigations of flares (see
e.g. Balona, 2015; Davenport, 2016) and superflares ((E$>$1033 erg), see e.g.
Maehara et al., 2012; Tu et al., 2020; Doyle et al., 2020; Tu et al., 2021;
Okamoto et al., 2021). Flare-frequency distributions from TESS or Kepler are
determined for more energetic flares, as usually such broadband photometric
observations are insensitive to low energetic flares, as those simply leave no
signature in a light-curve and are hidden in the noise.
Stellar CMEs have been detected so far mainly on dMe stars (e.g. Houdebine et
al., 1990; Guenther & Emerson, 1997; Vida et al., 2016) using the method of
Doppler shifted emission. This method uses the signature of plasma being
ejected from a star. The signature, either appearing in absorption or emission
is Doppler shifted by its projected velocity. This signature is often very
pronounced in Balmer lines similarly to erupting filaments/prominences on the
Sun. Optical spectroscopic monitoring programs to search for stellar CMEs
using the method of Doppler shifted emission/absorption are often focused on
dMe stars, as those are known to frequently flare and therefore also possibly
may host CMEs. dG stars reveal also frequent flares if their X-ray luminosity
is large, i.e. their activity level is high. However, for the detection of
flares and CMEs, only stronger events, compared to dMe stars, may be detected,
because of the higher continuum around H$\alpha$ on those stars, or in other
words a higher contrast is favorable for stellar CME detection using the
method of Doppler-shifted emission/absorption in the H$\alpha$ line.
Only recently Namekata et al. (2021) presented the detection of a blue wing
absorption in H$\alpha$ during a superflare on the young solar analogue EK
Dra, which the authors interpreted as an erupting filament. This event was
simultaneously observed by two telescopes. Inoue et al. (2023) report on the
detection of a high velocity blue-wing emission feature being interpreted by
the authors as prominence eruption on the RS CVn system 1355 Ori, consisting
of a K2-4 sub-giant and a G1 dwarf. Even more recently Namekata et al. (2024)
focus again on EK Dra and present this time two prominence eruptions from
which one has a projected bulk velocity being above the escape velocity of EK
Dra and reveals a simultaneously observed candidate of coronal dimming.
Using the method of Doppler shifted emission/absorption numerous candidate
events have been found especially on dMe stars (e.g. Fuhrmeister et al., 2018;
Vida et al., 2019). With this method only events with a projected bulk
velocity being greater the stars’ escape velocity can be treated as eruptive
events. It can be concluded that they are escaping from the star, as their
true velocities can be even higher. But in numerous studies events with much
smaller projected bulk velocities have been found which are much more
difficult to interpret, as those may also originate from flaring plasma
motions. To better interpret these spectral signatures investigating the Sun
seen as a star may help. Instruments doing solar 2D spectroscopy are rare. In
the 90ies of the last century Mees CCD provided H$\alpha$ 2D spectroscopy of a
cut-out of the solar-disk, in 2016 SMART/SDDI went into operation, doing full-
disk multi-filter measurements resulting in a full-disk H$\alpha$ profile, and
only recently the Chinese H$\alpha$ Solar Explorer (CHASE) providing full-disk
H$\alpha$ spectroscopy. Mees did H$\alpha$ 2D H$\alpha$ spectroscopy of solar
cycles 22 and 23, so representing the past, whereas SMART/SDDI operated during
the second half of cycle 24 and cycle 25 and CHASE beginning of cycle 25 up to
now and hopefully also in the future. Spectroscopic Sun-as-a-star observations
date back to the seventies of the last century (e.g. Livingston et al., 1981).
Spatial integrated investigations of flares and erupting filaments have been
presented by e.g. Den & Kornienko (1993); Ding et al. (2003); Ichimoto et al.
(2017) and only recently progress has been made especially to understand
Balmer line asymmetries related to erupting filaments/prominences (Namekata et
al., 2021; Leitzinger et al., 2021; Namekata et al., 2022b; Otsu et al., 2022;
Otsu & Asai, 2024; Ma et al., 2024), as the Sun is the only star where we can
actually see if a CME occurred in spatial and temporal vicinity to a
filament/prominence eruption.
Also solar CME signatures are used to be searched on other stars, such as
radio type II and IV bursts. Solar type II bursts are signatures of shock
waves which can be driven by CMEs. On the shock front electrons are
accelerated and radiation is emitted in dependence of the plasma frequency.
Moving type IV bursts are related to trapped electrons inside CME structures
originating from flare sites (see e.g. Gopalswamy, 2016). Several attempts
have been made to detect those radio signatures, at low frequencies (e.g.
Leitzinger et al., 2009; Boiko et al., 2012; Konovalenko et al., 2012; Crosley
et al., 2016, , in the MHz regime) and also at higher frequencies (e.g.
Crosley & Osten, 2018; Villadsen & Hallinan, 2019; Zic et al., 2020; Bloot et
al., 2024, in the GHz regime). In nearly all of those studies bursts have been
detected but so far no study claimed that a stellar analogue of a solar radio
type II burst was detected. An analogue of a solar type IV burst occurring on
Proxima Centauri has been presented by Zic et al. (2020). Only recently Mohan
et al. (2024) reported on the detection of a solar-like radio type IV burst on
the young and nearby dMe star AD Leo using 550-850 MHz observations from the
upgraded Giant Meterwave Radio Telescope (uGMRT). Mullan & Paudel (2019)
argued for radio quiet CMEs as the strong magnetic fields in dM stars would
require very large CME velocities to produce type II emission. Ó Fionnagáin et
al. (2022) investigate $\epsilon$ Eri theoretically and find that type II
bursts are visible above the ionospheric cut-off frequency sufficiently long
(about half an hour) if observed by LOFAR. Alvarado-Gómez et al. (2020)
investigate type II burst frequencies on Proxima Centauri and find that type
IIs may emit below the ionospheric cut-off frequency making it unaccessible to
ground-based radio facilities.
Continuous absorptions during flares at X-ray wavelengths have been observed
and interpreted being caused by plasma clouds obscuring the stellar flaring
region. Favata & Schmitt (1999) find during a flare on Algol an increased
hydrogen column density which decreased in time. This behaviour was
interpreted by the authors as cool, absorbing material in the line of sight,
reminiscent of an erupting filament (see Moschou et al., 2019, for a review on
this method).
Another solar CME signature is coronal dimmings (for a review on solar CME
detections see Webb & Howard, 2012) which is the sudden evacuation of coronal
plasma on the Sun being closely correlated to CMEs (e.g. Dissauer et al.,
2019). By establishing the Sun-as-a-star signature of coronal dimmings on the
Sun, this method was applied to stars in Veronig et al. (2021) where around 20
potential CME candidates have been detected on late-type main-sequence stars.
There are also other signatures which authors related to be caused by CMEs,
such as e.g. Bond et al. (2001), who found transient absorption features in a
pre-cataclysmic binary system containing a white dwarf and a K-dwarf which the
authors interpreted as CMEs from the K-dwarf crossing the line-of-sight in
front of the white dwarf. Şenavcı et al. (2018) investigate the RS CVn
eclipsing binary SV Cam (F9V+K4V) and identify excess absorption which the
authors interpret as cool plasma (filaments/prominences) obscuring the primary
component.
For reviews on stellar CMEs see Moschou et al. (2019), Leitzinger & Odert
(2022), Osten (2023), Tian et al. (2023).
For the detection of stellar CMEs the approaches are manifold as can already
be seen from the methodologies mentioned above. For the method of Doppler
shifted emission/absorption the observing strategies are manifold as well. One
can either observe single stars or one uses multi-object spectroscopic devices
which are able to observe several stars at the same time spectroscopically
(see e.g. Guenther & Emerson, 1997; Leitzinger et al., 2014; Korhonen et al.,
2017). Both approaches have their advantages but also their drawbacks. Single
star observations are of course very time consuming, because when one focuses
not only on one star then observing time increases by factors in the order of
the number of stars one wants to observe. For multi-object observations one
catches in one exposure several stars but the target selection is of course
limited. Existing multi-object spectrographs have very differently sized
fields-of-views (FOVs), from square arcmin to square degrees. But this means
always one has to select open clusters or associations to fill the FoV of the
instrument with a sufficiently large number of targets of similar age.
Moreover, multi-object spectrographs are mainly available on large telescopes
where usually the access is restricted to member institutions or telescope
time is distributed via a competition process. Experience has shown that
observing proposals with the aim of monitoring stars are not likely granted
with observing time as such proposals would block a telescope over a long
period. Therefore another approach needs to be chosen for monitoring of solar-
analogues. At the authors home institution a small-sized telescope is
available. The 0.5m telescope is equipped with a slit spectrograph and a back-
end device which was optimized for this purpose. The minimum requested signal-
to-noise (S/N) for the detection of CMEs within 10-15 minutes exposures, to
properly resolve CME evolution, is $\sim$20\. This restricts the target
selection to brighter stars ($<$7 mag). There is a small number of young solar
analogues which fulfill this criterion (see section 2.2). As the telescope
time can be attributed up to 75% of the total observing time we
spectroscopically monitored solar analogues for three years to search for
optical signatures of flares and erupting filaments/prominences in H$\alpha$
using the method of Doppler-shifted emission/absorption. .

##  2 Observations, data, and targets

###  2.1 Observatory Lustbühel Graz (OLG)

The OLG is situated at the border of the city of Graz, Austria. It contains an
astronomy part belonging to the Institute of Physics, Department for
Astrophysics and Geophysics of the University of Graz, a satellite geodesy
part, belonging to the Space Research Institute of the Austrian Academy of
Sciences, and a telecommunications part belonging to the University of
Technology, Graz, Austria. The building hosts two domes, one belonging to the
astronomy part, hosting a 0.5m reflector manufactured by the Astro Systeme
Austria (ASA) company and the other to the satellite geodesy part. The ASA
0.5m is a german mount telescope with a focal length of 4.5m. Separated from
the building exists a tower with a third dome hosting the Ballistische Mess
Kammer (BMK), a rare astrograph manufactured by ZEISS from the nineteen
seventies, with a focal length of 0.75m yielding an immensely large angular
FoV of 19.3 degrees.

####  2.1.1 Instruments

The ASA 0.5m telescope is equipped with four instruments mounted on a custom
manufactured so-called multi-port, a device mounted in Cassegrain focus with a
rotating mirror feeding light to the instruments. Mounted to the multi-port
are two imaging cameras with broad and narrow-band filters, one video camera
for asteroid occultation measurements, and a spectrograph with a dedicated CCD
camera.
The spectrograph is a Littrow High Resolution Spectrograph (LHIRES), an off-
the-shelf slit spectrograph with adjustable wavelength range and removable
gratings. The spectrograph back-end is an Apogee Alta F47 back-illuminated CCD
camera reaching a quantum efficiency of $>$90% around H$\alpha$. For
spectroscopic monitoring we select the medium/low resolution grism with 600
groves/mm corresponding to a spectral resolving power of R=2700 which
corresponds to $\sim$100 km s-1 at H$\alpha$. This is a resolution comparable
to the now already decommissioned VIsible Multi Object Spectrograph (VIMOS) at
the Very Large Telescope (VLT) of the European Southern Observatory (ESO) with
its orange grism which we have used in the past (see Leitzinger et al., 2014).
The used grism yields a spectral range of $\sim$5800 .. 6800Å.

Table 1: Target stars of the study, their characteristics (taken from Güdel, 2007), number of observed spectra, and flare rates. The XUV flare rates have been estimated from the X-ray luminosities using the power-law from Audard et al. (2000), the cumulative H$\alpha$ flare rates have been estimated from the XUV flare rates based on Leitzinger et al. (2020), and the TESS flare rates have been determined by eye from the light curves. HN Peg and $\kappa^{1}$ Cet did not show any flares in their TESS light curves therefore we give here upper limits only. name | age | log Lx | V | Prot | number of spectra | XUV flare rate | H$\alpha$ flare rate | TESS flare rate
---|---|---|---|---|---|---|---|---
| [Gyr] | erg s-1 | [mag] | day |  | N(>1032erg) day-1 | N(>1032erg) day-1 | day-1
EK Dra | 0.1 | 29.93 | 7.6 | 2.75 | 6722 | 52 | 2.81 | 0.54
HN Peg | 0.3 | 29.12 | 5.9 | 4.86 | 13433 | 9 | 0.48 |  $<$0.11
$\pi^{1}$ UMa | 0.3 | 29.10 | 5.6 | 4.68 | 5112 | 8 | 0.46 | 0.06
$\chi^{1}$ Ori | 0.3 | 28.99 | 4.4 | 5.08 | 6856 | 7 | 0.36 | 0.04
$\kappa^{1}$ Cet | 0.75 | 28.79 | 4.9 | 9.2 | 648 | 4 | 0.23 |  $<$0.06

####  2.1.2 Observing

Observing at OLG can be done almost automatically. Co-author R. Greimel
developed a python based software with user interface which combines all the
necessary software packages/tools needed to operate the telescope, dome, and
instruments. For the spectrograph a tool was developed which tracks the star
in the image of the guide camera hereby keeping the spectrograph slit on the
target star. This enables that the observer needs only to start the system,
take calibration frames, select the target of the night, and start observing.
The software system also includes an auto-close option which closes the dome,
parks the telescope, and disconnects all the software tools from the python
based software. To further ease the observer’s work the whole system can be
operated remotely and therefore the observer does not need to be at the
observatory to start the system. Furthermore OLG has a weather station
including all-sky camera, wind, temperature, and humidity measurements, all of
that is necessary, especially for remote observing.

###  2.2 Target stars

As already mentioned in the introduction, according to the aperture of the
telescope and the location of the observatory, we are restricted to bright
solar analogues. We decided to select only a handful of solar analogues to
dedicate sufficient observing time to each target to enhance the chance to
detect signatures of flares and CMEs on those stars. We used the list of the
“Sun in time” sample (Güdel, 2007) to compile a target list. In table 1 we
list the target stars of our study together with characteristics and the
number of recorded spectra for each target star.
EK Dra: EK Dra is a $\sim$100 Myr solar analogue star of spectral type dG0
with a logLx being $\sim$100 times larger than for the Sun (see e.g. Güdel,
2007; Şenavcı et al., 2021) and rotating 10 times faster than the Sun (König
et al., 2005). EK Dra has a wide companion with a period of $\sim$45 years
having roughly half the mass of the primary (König et al., 2005).
HN Peg: HN Peg is $\sim$300 Myr solar analogue of spectral type dG0 with a
logLx being $\sim$50 times larger than the solar value (see e.g. Güdel, 2007)
and rotating $\sim$6 times faster than the Sun. HN Peg has a T-dwarf companion
with a separation of $\sim$800AU (Luhman et al., 2007).
$\pi^{1}$ UMa: $\pi^{1}$ UMa has also an age of 300 Myr and a logLx being 40
times larger than that of the Sun (see e.g. Güdel, 2007) similar to HN Peg and
$\chi^{1}$ Ori. Also the rotation period is comparable to the ones of HN Peg
and $\chi^{1}$ Ori making it rotate $\sim$6 times faster than the Sun.
$\pi^{1}$ UMa is a single star with no known companion.

Figure 1: TESS light curves of the target stars from top to bottom for EK Dra
(sector 21), HN Peg (sector 55), $\chi$1 Ori (sector 44), $\pi$1 UMa (sector
60), and $\kappa$1 Cet (sector 31). Arrows mark flares in the light curves.

$\chi^{1}$ Ori: $\chi^{1}$ Ori is also a $\sim$300Myr solar analogue of
spectral type dG1 with a logLx being $\sim$30 times larger than that of the
Sun (see e.g. Güdel, 2007) and rotating $\sim$5 times faster than the Sun,
similar to HN Peg. $\chi^{1}$ Ori has a companion of spectral type M with a
period of $\sim$14 years and a mass of one seventh of the mass of the primary.
$\kappa^{1}$ Cet: $\kappa^{1}$ Cet is the oldest star in the sample with its
750 Myr. Its spectral type is dG5 and its logLx is $\sim$20 times larger than
the solar value (see e.g. Güdel, 2007). $\kappa^{1}$ Cet spins $\sim$3 times
faster than the Sun. Companions of $\kappa^{1}$ Cet have been searched but to
date none has been confirmed.

###  2.3 Other data and data reduction

For deducing TESS flare rates we use all available TESS 2 minute light curves
for the stars in our target sample (see Fig. 1 for examples). For the
identification of flares we use the already extracted light-curves available
from TESS. For EK Dra we use sectors 14, 15, 16, 21, 22, 23, 41, 48, 49, 50,
for HN Peg sector 55, for $\chi^{1}$ Ori sectors 43, 44, 45, for $\pi^{1}$ UMa
sectors 20, 47, 60, and for $\kappa^{1}$ Cet sectors 4, 31. As the number of
light curves to be searched for flares is small (19) we analysed the light
curves by eye. EK Dra shows in every TESS light curve flares. $\pi^{1}$ UMa
reveals only one flare, in one out of three TESS sectors. $\chi^{1}$ Ori
reveals only three flares in one out of three TESS sectors and HN Peg and
$\kappa^{1}$ Cet reveal no flares in their TESS light curves.
For a comparative analysis of solar activity phenomena (see section 3.2) we
utilize data from Mees CCD imaging spectrograph (MCCD, Penn et al., 1991) at
Mees Solar Observatory (MSO). These data are two dimensional imaging
spectroscopic data and are available as level 0.5 and/or level 1.0 data. Level
0.5 data are raw frames with a fits header whereas level 1.0 data are
nominally processed. For the solar analysis we used level 1.0 data with a
modified implementation of the default flat-fielding code (priv. comm. KD
Leka).
Optical spectroscopic observations for the study started in June 2018 and
lasted until September 2021, with a break between November 2018 and March 2019
where the camera had to be sent back for electronic adjustment. For
spectroscopic observations we used only half of the chip to minimize read-out
time and save hard drive memory. The thereby reduced number of pixels was
still sufficient to determine the background properly. The usual exposure time
was set to three minutes as this ensures a proper temporal resolution and in
case of low signal, exposures can be added without increasing noise
significantly. The data were reduced with the Imaging and Reduction Analysis
Facility (IRAF) using the package $specred$ and background was subtracted
using the task $apscatter$. The spectra were wavelength calibrated and
normalized before analysis. The spectra show a drift per night due to the fact
that the spectrograph is Cassegrain mounted and therefore flexure is an issue.
The drift in wavelength needs to be removed which we have done using cross
correlation of the actual spectra with a reference spectrum per night. This
yielded one dimensional spectra which were then further processed using the
Interactive Data Language (IDL). In total we recorded 32771 spectra for the
investigation of activity of solar analogues (see table 1).

Figure 2: Examples of OLG spectra of the target stars from top to bottom for
EK Dra, HN Peg, $\chi$1 Ori, $\pi$1 UMa, and $\kappa$1 Cet.

##  3 Results

###  3.1 Stellar observations

We analyzed all OLG spectra in terms of variability in H$\alpha$ by computing
the residual spectra by subtracting a quiescent spectrum from the time series.
The quiescent spectrum is represented by the mean of all spectra from one time
series, i.e. one night, after rejecting noisy spectra. In these residuals we
did not see signatures of flares or CMEs. As a second step, we then binned the
spectra in time up to 30 minutes to increase S/N.
The typical S/N of the stellar spectral sample vary from star to star. For the
nominal integration times (3 min) we derive a S/N distribution for EK Dra
peaking at a value of 60, for HN Peg peaking at a value of 80, for
$\pi^{1}$UMa peaking at the same value, for $\chi^{1}$Ori peaking at a value
of 110, and for $\kappa^{1}$ Cet peaking at a value of 90. The S/N values
behave accordingly with the stars magnitudes, which is reasonable. The S/N
values were calculated with the algorithm presented in Stoehr et al. (2008)
which represents a general way to derive S/N from one-dimensional spectra.
Temporal binning of up to 10 spectra yields S/N distributions with peak values
being 30-40% larger than that derived from the single spectra, but with
maximum values being a factor 2 larger than the S/N distributions derived from
the single spectra. The algorithm from Stoehr et al. (2008) seems to be more
conservative in deriving S/N values when compared to the rough estimation of
spectral noise for bright stars being calculated as the square root of the
stellar flux.
Of course, with temporal binning the number of temporally binned spectra
decreases, so that we end up with 20 spectra per night at maximum,
corresponding to one full winter night of $\sim$10 hours. Also here we
produced residual spectra. We produced a catalogue of spectra and residual
spectra which we then analysed by eye. Even if the amount of nights is quite
large, signatures lying above the noise are easily visible, and therefore a
search for signatures is feasible by eye. This revealed signatures which were
invisible in the original unbinned three minute exposures. We did not see any
signatures of flares and CMEs in the target stars except for EK Dra being the

Figure 3: Upper panel: TESS lightcurve coinciding in time with the residual
spectral time series shown in the lower panel. Differently grey-shaded areas
denote the residual spectra in the lower panel. The indexing corresponds to
the indexing in the lower panel. Lower panel: residual spectral time series
(from top to bottom) of EK Dra from the night of the 5th April 2020. Grey
shaded areas denote the 1-$\sigma$ uncertainty of the residuals. Red solid
lines are gaussian fits to the absorption visible from residuals 1-5. Table 2:
Detected flares/CMEs and their characteristics from the present study (all
from EK Dra). Calculation of the flare energies is described in the text.

observing night | flare | CME | starting time H$\alpha$/TESS | flare duration H$\alpha$/TESS | flare energy H$\alpha$/TESS | H$\alpha$ asymmetry
---|---|---|---|---|---|---
|  |  | [BJD-2450000] | [min] | [erg] |
19/02/2020 | yes | no | 8899.5174/8899.5322 | 150/54 | 2.3$\times$1032/3.5$\times$1033 | red
04/04/2020 | yes | no | 8944.5918/8944.5791 | 60/48 | 6.8$\times$1031/7.0$\times$1033 | red
05/04/20201 | no | yes | 8945.2529/- | -/- | -/- | blue
14/04/2020 | yes | no | 8954.2819/8954.3164 | 240/60 | 8.0$\times$1032/1.7$\times$1033 | red
22/04/2020 | yes | no | 8962.5232/- | 180/- | 6.9$\times$1032/- | red

(1)We know from the TESS light curve and Namekata et al. (2021) that there was
a flare prior to event 05/04/2020, but we did not detect the flare by the
spectroscopic monitoring, as the spectroscopic time series that night started
after the flare.

youngest and most active but also the faintest solar analogue in the sample,
and therefore the star with the lowest S/N in the spectra.
We observed EK Dra for $\sim$336 hours which is a similar amount of time as
for $\chi^{1}$Ori or for $\pi^{1}$UMa. TESS observed EK Dra in sectors 14-16,
21-23, 41, and 48-50, and we observed coordinated with TESS in sectors 21-23.
The TESS light curves of sectors 21-23 are shown in the appendix in Fig. 10.
We have indicated the times where we have observed spectroscopically by blue
shaded areas. What becomes obvious is that we missed all larger flares.
However, we caught some of the smaller flares, namely three in TESS sector 23
named after the nights in which they occurred, event 19/02/2020 (see left
panel of Fig. 12), event 04/04/2020 and 14/04/2020 (see Fig. 11). Another
flare was detected after TESS sector 23, so not coordinated with TESS, named
event 22/04/2020 (see right panel of Fig. 12). An absorption signature,
reminiscent to a solar eruptive filament, named event 05/04/2020 in TESS
sector 23 was seen as well. All of the TESS flares were superflares which can
be seen from the estimated TESS flare energies given in Table 2. In Table 2 we
list starting times of the detected flares and the eruptive filament on EK
Dra, flare durations, and flare energies derived from TESS and/or H$\alpha$
data. The TESS flare durations were calculated according to the pre-flare
level in the light curves and the time when the light curves reach pre-flare
levels again. The TESS flare energies are calculated by extracting the TESS
flux (8.12$\times$10-9erg/s/cm2/Å) from the TESS magnitude of EK Dra
(Tmag=7.04, Paegert et al., 2022), using the distance to EK Dra (d=34.45 pc,
Petit et al., 2014) and then integrating the residual light curves (in flux
units) over the flare durations. We also list the H$\alpha$ flare energies
which we derive from the residual spectral time series by applying the quiet
continuum flux around H$\alpha$ derived from GAIA DR3 spectra of EK Dra
(2.75$\times$10-12erg/s/cm2/Å) which are available as flux calibrated spectra
(De Angeli et al., 2023). By involving the distance to EK Dra and integrating
over the H$\alpha$ flare durations we are able to obtain H$\alpha$ flare
energies. These energies are lower limits as we do not take into account the
continuum increases during the H$\alpha$ flares However, for our flares which
have been observed by TESS simultaneously, we know that the continuum
variation is small (see description of single events below) and therefore
negligible. This does not account for flare 22/04/2020 as this one was not
observed simultaneously by TESS.
In the following we describe the events from Table 2 in more detail.
Event in night 19/02/2020 in TESS sector 22: TESS captured a relatively small
flare with an amplitude of $\sim$0.3% above background which lasted for about
an hour. The impulsive flare phase started around BJD(-2450000) 8899.5174 and
after $\sim$5 min the flare had reached its peak. The spectroscopic
observations started that night one binned spectrum prior to the flare as
detected by TESS (BJD(-2450000) 8899.4983, see Fig. 12). Already the second
spectrum covers the impulsive TESS flare phase and spectra 3 and 4 cover the
decay phase of the TESS flare (see TESS light curve, upper panel of Fig. 12).
Taking a look at the residual spectra in the left lower panel of Fig. 12
reveals that in H$\alpha$ the flare signature can be seen beyond the decaying
phase of the TESS flare deduced from TESS broadband photometry until spectrum
6. The residual spectra show a peak being red shifted by one wavelength bin
(similarly as event 14/04/2020 in TESS sector 23).
Event in night of 04/04/2020 in TESS sector 23: This TESS flare starting at
BJD(-2450000) 8944.575 and reaching its peak in $\sim$10 min is also a
relatively weak flare with an amplitude of $\sim$0.4%. The flare was
spectroscopically captured at the end of the night. The flare peak is covered
by spectrum no. 13 which does not reveal a significant increase in flux at
H$\alpha$. In the successive spectrum a distinct, broad (FWHM=10.4Å), and red-
shifted ($\Delta\lambda$=2.1Å) signature is visible which gains flux in the
following residuum. All this happens in the decaying tail of the TESS flare.
As the light curve returns to quiescent levels the red-shifted asymmetry
vanishes in the last residual spectrum.
Event in night of 14/04/2020 in TESS sector 23: This TESS flare starts at
BJD(-2450000) 8954.316 and reaches its peak around 15 min later. After 30 more
minutes the flare has reached quiescent levels again. The peak of the flare is
only $<$0.2% above the background. The residual spectra reveal, that there is
already excess flux in the residual spectrum prior the residual spectrum
covering the impulsive phase. The excess flux in the residual spectra peaks
during the TESS flare decay phase (cf. right lower panel of Fig. 11). We still
see some excess emission in the H$\alpha$ residual spectra although the
broadband flare has already ended. Until residual spectrum 9 (cf. right upper
panel of Fig. 11) we see excess emission and this is five residual spectra or
$\sim$2.5 hr later. Also here (similarly to event 19/02/2020 in TESS sector
22) we see a red-shift of the excess emission peak by one wavelength bin.
Event in night 22/04/2020: This flare was not observed coordinated with TESS.
From the residual spectra (BJD(-2450000) 8962.5232 .. 8962.6082; see Fig. 12)
we see that there are broad emission features (FWHM=6.8-9.7Å) in H$\alpha$.
These flare signatures are detected in five spectra indicating a total
duration of $\sim$2.5 hr. As time evolves the signatures move to the red,
shifted at maximum by 3.2Å which corresponds to $\sim$150 km s-1.
Event in night 05/04/2020 in TESS sector 23: TESS detected a flare starting at
BJD(-2450000) 8945.238 which reached its peak after seven minutes. The decay
of the flare lasts for $\sim$15 min. Our coordinated spectroscopy started at
the time when the star had already returned to its quiescent level in the TESS
light curve. We recognized a significant absorption being visible already in
the first spectrum lasting for five spectra (see Fig 3), i.e. 2.5 hr. This
broad absorption feature (FWHM=7.5 .. 14.9Å) moved from the blue to the red,
crossing the H$\alpha$ core and then moving back to the H$\alpha$ core
wavelength. This event is reminiscent to a solar filament eruption, for a
detailed discussion see section 4.
As we have detected no flares and/or CMEs for the remaining stars in the
sample we may derive upper limits only. As in Leitzinger et al. (2020) we use
the relation -ln(1-0.95)/tobs from Gehrels (1986) to derive upper limits of
CME rates for the target stars with no detections. Doing so for our
observations yields an upper limit for the CME/flare rate of HN Peg of 0.11
day-1 ($\sim$1 in 5 days), of $\chi^{1}$ Ori of 0.21 day-1 ($\sim$1 in 5
days), of $\pi^{1}$ UMa of 0.28 day-1 ($\sim$1 in 4 days), and of $\kappa^{1}$
Cet of 2.22 day-1. For EK Dra we have detected only one event which resembles
a filament eruption, which showed a maximum projected velocity of the bulk
material being close to the stars’ escape velocity (cf. Namekata et al.,
2021), and therefore has likely evolved into a CME. The observed CME rate
yields 0.07${}^{+1.61}_{-0.058}$ day-1, which is a lower limit as the true CME
rate must be higher as what we are able to detect with our observational
setup. EK Dra is also the only star in the sample where we have detected
flares in H$\alpha$, therefore the observed flare rate yields
0.29${}^{+0.229}_{-0.139}$ H$\alpha$ flares day-1. In section 4 we relate
those values to the upper limits of our target stars in Leitzinger et al.
(2020) as well as

Figure 4: First column: Spatially integrated (for integration area see
white/black rectangles in columns 2-4) H$\alpha$ spectra (black: current, red:
quiescent). Second column: H$\alpha$ core image (at 6562.3Å). Third column:
H$\alpha$ blue wing image (at 6560.2Å). Fourth column: H$\alpha$ red wing
image (at 6565.4Å). Every row (from top to bottom) corresponds to the
numbering in Fig. 5 (dot symbols).

relate the upper limits to expected CME rates (Odert et al., 2017, 2020). We
will also discuss the expected flare rates as well as the recent detection of
a filament eruption on EK Dra (Namekata et al., 2021).

###  3.2 Solar observations

To evaluate the detectability of flares and CMEs on our target stars we
construct Sun-as-a-star H$\alpha$ spectra of erupting filaments and flares
(see Fig. 4), utilizing observations from MCCD at MSO which has observed the
Sun between 1991 and 2007, covering thereby two maxima of solar cycles 22 and
23. In its flaring mode MCCD has a temporal resolution of $\sim$16 s, a
spectral resolution of 0.375Å and a FoV of 3.8$\times$4.8 arcmin.

Figure 5: Solar time series of the H$\alpha$ core intensity of the 15/07/2002
event (black solid line) overplotted with the corresponding GOES X-ray light
curve (grey solid line). Numbers correspond to times when flares, filament
eruptions, backflowing material, and a Halo CME occurred (see text).

The data archive111https://www.cora.nwra.com/MCCD/ is freely accessible and
contains numerous flares and filament/prominence eruptions. Therefore we
selected all X-, M-, and C-class flares from
GOES222https://www.ngdc.noaa.gov/stp/space-weather/solar-data/solar-
features/solar-flares/x-rays/goes/xrs/, and crossmatched the peak times of the
flares with the observing times at MCCD/MSO. Thereby we identified $\sim$40
filament/prominence eruptions. The most significant event is the filament
eruption and X-class flare from 15/07/2002, which is a well investigated event
(see e.g. Liu et al., 2003; Gary & Moore, 2004; Li et al., 2005). We selected
this event to investigate the detectability of flares and CMEs on solar
analogues as well as to determine the characteristics of active regions
causing such eruptions. We construct Sun-as-a-star spectra by summing up all
pixels belonging to the rectangle, which we have defined to include the
flaring area as well as the eruptive filaments. For the strong filament
eruption shown in Fig. 4 in the sixth row (from the top) one can see that our
defined rectangle does not cover the complete area of the filament, which does
not significantly affect the spatially integrated H$\alpha$ spectrum. The
15/07/2002 event started with an X3 flare at 20:03UT on 15th July 2002 near
disk center in active region NOAA 10030 followed by a filament eruption.
Around half an hour later a second weaker flare peaked, occurring during the
decaying tail of the X-class flare, with another filament eruption. A Halo CME
was seen at 20:30UT and at 21:06 a partial Halo CME occurred.
In Fig. 5 we show the H$\alpha$ core intensity of the MCCD observations of the
event from 15/07/2002 (black solid line) overplotted with the corresponding
GOES X-ray light curve (grey solid line). We indicate seven different times in
the light curve (dot symbols) for which we show the spatially integrated
H$\alpha$ spectrum, together with the H$\alpha$ core, blue-(at $\sim$656.02
nm) and red-wing $\sim$656.54 nm) image in Fig. 4. In addition we show also
the position of the peak of the corresponding GOES X-ray light curve on the
H$\alpha$ light curve as well as the occurrence of the associated Halo CME.
Number 1 in Fig. 5 corresponds to the H$\alpha$ pre-flare stage, number 2
indicates the onset of the first

Figure 6: Time series (from top to bottom) of spectral residual flux extracted
from the active region in Fig.4 (white rectangle) of the solar event from
15-07-2002. The residual spectra have been scaled with a factor resulting in
an area scaling of 18 %, which is the scaling with which the spectral
signatures of the eruptive filament are detectable exceeding 1-$\sigma$. The
grey shaded area corresponds to the median 1-$\sigma$ uncertainties from the
spectral residuals (original 3 minute spectra) of all events detected on EK
Dra (see Table. 2). Each residual spectrum is tagged with a time which
corresponds to the times shown in Fig. 5, as well as selected residual spectra
with a number which corresponds to the indexing in Fig.5. Figure 7: Upper
panel: Cut-out of a normalized HN Peg H$\alpha$ spectrum (black solid line).
The blue area marks the standard deviation. The green solid line represents
the spatially integrated solar MCCD spectrum of one time step (t=20.6 hour in
Fig.8) of the erupting solar filament from 15-07-2002 (see lower panel)
superimposed on the HN Peg spectrum. From these superimposed spectra (green
solid line) we determined the EW time series shown in Fig.8. Lower panel:
Solar spatially integrated residual H$\alpha$ spectrum of one time step of the
erupting solar filament from 15-07-2002. The applied area scaling is here 24%.

Figure 8: Upper panel: Time series of equivalent widths of the solar H$\alpha$
residual profile of the 15/07/2002 event superimposed on the H$\alpha$ profile
of HN Peg of the original 3 min exposures. The depression starting at 20:05 UT
corresponds to the X3 flare (a flare causes a decrease in EW and an increase
in flux), the increase at 20:35 UT corresponds to the very pronounced filament
eruption corresponding to number 6 in Fig. 5. The maximum of the EW time
series corresponds to both absorption signatures, blue and red, i.e. outwards
moving and back flowing material (at 20:10 UT). The different colored curves
correspond to time series of equivalent widths scaled with factors 40, 50,
100, and 150, which in turn corresponds to fractional active area 24, 30, 60,
and 90%. The black solid line is the original non-scaled equivalent width time
series. The grey shaded area corresponds to the $\pm$ 1-$\sigma$ error of the
spectrum of HN Peg. Lower panel: same as upper panel but with half hour
averages, which are the averages which have revealed the signature of the
erupting filament on EK Dra. As once can see, the signature of the erupting
filament has vanished completely as the half hour average bin includes the
times when the second flare and the filament eruption happened.

filament eruption, number 3 indicates the peak of the first flare (GOES X3),
number 4 indicates the peak of the GOES X-ray light curve, number 5 indicates
the onset of the second filament eruption, number 6 indicates the second
weaker flare, number 7 indicates the onset of the third and very pronounced
filament eruption, number 8 indicates the occurrence of the associated Halo
CME and number 9 indicates the first occurrence of back-flowing filament
material.
As one can see in Fig. 4, the flares and the filament eruptions leave distinct
signatures in the spatially integrated spectra (first column). A flare causes
spectral line core enhancements (see Fig. 4 third row), an erupting filament
causes, when moving towards the observer, a blue-shifted absorption (see Fig.
4 fourth and sixth row), and back-flowing filament material causes red-shifted
absorptions (see Fig. 4 last row). The strength of the signatures is strongly
dependent on the size of the integration area (white and black, respectively,
rectangle in the H$\alpha$ core, blue-, and red-wing images shown in columns
2-4).
To evaluate if we could detect a typical solar flare and filament eruption in
our stellar spectra we use the event from 15/07/2002 in two approaches. For
both approaches we calculate the normalized residual flux as fnorm =
Aactive/A${}_{star}\times$(Iactive/Iquiet \- 1), where Aactive is the area of
the active region, Astar is the area of the stellar disk, Iactive is the
spatially integrated solar intensity of the active region, and Iquiet is the
spatially integrated solar intensity of the quiet Sun (determined from the
same region but from the first spectral image of the time series).
In the first approach we aim to see with which area scaling we may overcome
the 1-$\sigma$ uncertainties of the residual fluxes of the solar event from
15-07-2002 , similarly to the stellar case shown in Fig.3. In Fig.6 we plot
the residual fluxes obtained from the solar event from 15-07-2002 as 3-minute
averages (nominal integration time of the stellar spectroscopic monitoring)
scaled to an active region area of 18 % of the solar disk. With such an area
scaling the spectral residual signatures of the eruptive filament exceed
1-$\sigma$. Without any area scaling the residual spectral signatures remain
hidden in the noise.
As a second approach we superimpose the solar residual event spectrum onto the
spectrum obtained for one of our solar analogue target stars and determine the
H$\alpha$ equivalent widths. To do so we adapt the wavelength grid of the MCCD
observations to the OLG observations and apply the area scaling as described
above for approach 1. Here, the term fnorm which is the solar residual
H$\alpha$ spectrum, is then simply added to the stellar spectrum (see Fig.7).
Applying the above formalism to the stellar OLG spectra reveals that the solar
event from 15/07/2002 would have been invisible if occurring for instance on
HN Peg, simply because the field of view of MCCD corresponds to $\sim$2.3% and
the integration area (white and black, respectively, rectangle in the
H$\alpha$ core, blue-, and red-wing images shown in columns 2-4 of Fig. 4)
only to $\sim$0.6% of the solar disk, and this is way to small to be seen with
respect to the solar disk in a Sun-as-a-star spectrum. Subsequently, we
generate in Fig. 8 a time series of equivalent widths (determined from the
full spectral coverage of MCCD, see Fig. 7) of the solar residual H$\alpha$
spectra of the 15/07/2002 event superimposed on the spectrum of HN Peg. The
integration area covering the solar active region is, as mentioned above, too
small to make the signature visible in a Sun-as-a-star spectrum. Therefore we
scale the active region area by various factors as shown in Fig. 8. The grey
shaded area denotes the EW error which we estimated using the relation given
in Nissen & Gustafsson (2018). As one can see that with a scaling factor of 40
which corresponds to a fractional area of 24% the filament signature is
sufficiently strong to be detected exceeding 1-$\sigma$. The solar H$\alpha$
flare seen as a depression in Fig 8 is exceeding 1-$\sigma$ with a fractional
area of $\sim$12%, i.e. it is much easier visible than a filament eruption. As
we have detected activity on EK Dra in the OLG data with temporal binning up
to 30 minutes we also binned the time series of H$\alpha$ equivalent widths to
see how that affects the evolution and to mimic stellar observations with
longer exposure times (see lower panel of Fig. 8). For a more detailed
discussion of this aspect see section 4.

##  4 Discussion

In the discussion, we relate the obtained upper limits of the CME rates of the
target stars to results obtained in Leitzinger et al. (2020), to the expected
observable CME rates from Odert et al. (2020), and to the recent detection of
a filament eruption presented in Namekata et al. (2021). We moreover discuss
the solar analysis in comparison with recent similar attempts from literature.
Stellar analysis: In Leitzinger et al. (2020) we followed a similar approach
as here but with archival data which usually have not the desired parameters
such as S/N and length of the time series. In Leitzinger et al. (2020) we
analyzed a much higher number of stars, among them were also the target stars
of the present study. As the analyzed time series in Leitzinger et al. (2020)
were shorter and no signatures of stellar CMEs could be detected, the present
observations constrain the upper limits of CME rates of the target stars given
in Leitzinger et al. (2020). For HN Peg we find an upper limit for the CME
rate of 0.11 day-1 which improves the prior limit (2) by a factor of 18, for
$\pi^{1}$ UMa we find an upper limit of 0.28 day-1 which constrains the prior
found limit (2.37) by a factor of $\sim$9, for $\chi^{1}$ Ori we find an upper
limit of 0.21 day-1 which constrains the prior found limit (4.28) by a factor
of $\sim$20, for $\kappa^{1}$ Cet we find an upper limit of 2.2 day-1 which
constrains the prior found limit (16.42) by a factor of $\sim$8, and finally
for EK Dra we find a detected CME rate of 0.07${}^{0.161}_{-0.058}$ day-1.
These upper limits are to be seen only with respect to the S/N of the spectra,
and therefore with CMEs of a certain mass. To relate upper limits of CME rates
of stars to the mass of CMEs modelling is necessary. In Leitzinger et al.
(2020) we have therefore utilized the semi-empirical model from Odert et al.
(2017, 2020) to obtain expected observable CME rates which we then compared to
the observed ones. The S/N of the spectra determines which CMEs with which
masses can be detected. The mean S/N of the observations used in Leitzinger et
al. (2020) for the target stars of the present study had a mean of 293,
roughly a factor 3 higher than the observations used in the present study.
Therefore for the Leitzinger et al. (2020) study CMEs with masses $\geq$
1.5$\times$1016g could have been detected whereas for the present study CMEs
with masses $\geq$5$\times$1016g could have been detected.
The semi-empirical modelling in Leitzinger et al. (2020) was done using the
code from Odert et al. (2020) for a fixed parameter set of optical thickness
in the H$\alpha$ line $\tau$=10, flare power law index $\alpha$=1.8, and
hydrogen column density Nh=1020 cm-2. Doing the same for the present
observations still the expected observable rates lie above the observed upper
limits, meaning that the expected observable rates overestimate the CME rates
of those stars. Varying $\tau$ (1 … 10), $\alpha$ (1.8 … 2.5), and Nh (1020 …
1021), yields observable expected CME rates lying above the observed upper
limits except the parameter combination Nh=1021 cm-2, $\tau$=10. This
indicates that plasma with an increased hydrogen column density together with
a high optical thickness agrees with the observed upper limits.
In contrast to the simple model above, hydrogen column density and optical
thickness are not independent and therefore this can be implemented in the
calculations for a self-consistent solution with respect to these two
quantities. Using fixed values for prominence parameters temperature and gas
pressure (Pgas=0.2 dyn cm-2, T=10000K) and the relationships presented in
Heinzel (2015) we self-consistently compute Nh and $\tau$. However, with this
modified code the observable expected CME rates overcome the observed upper
limits, meaning that with solar parameters the observable expected CME rates
are overestimated. We want to note, that already the intrinsic CME rates for
active stars are likely overestimated (Odert et al., 2017).
Moreover the observable expected CME rate calculations consider a
detectability within a 1-$\sigma$ uncertainty, leading to a minimum detectable
mass of 5$\times$1016g for the OLG observations. Considering a 3-$\sigma$
error range, then the minimum detectable mass increases to 1.5$\times$1017g,
which is in the order of the most massive solar CMEs. This finding coincides
with our expectations that with our observational setup we are able only do
detect signatures caused by massive stellar CMEs.
Only recently, a filament eruption seen in H$\alpha$ on the solar analog EK
Dra was presented by Namekata et al. (2021). Coordinated observations with
TESS were performed and one TESS flare (BJD(-2450000) 8945.24), was
accompanied by a filament eruption. As can be seen from Fig. 3 we also
observed in that period but we did not cover the flare. However, we started
observing shortly after this flare and we also detected the signature of an
erupting filament, but with lower temporal resolution as Namekata et al.
(2021), as these authors used larger observing facilities. In Fig. 1e of their
study they present the residual spectra of the event. The residual spectra
tagged with time ranges “60-80”, “80-100” and “100-120” are the ones which we
have captured also, although we have captured only 10 minutes of their “60-80”
average. However, when comparing our 30 minute averages with their 20 minute
averages, the shape of the absorptions differ to some degree, as the averages
are overlapping by 10 minutes. The absorption depth of our first two spectra
(see Fig. 3) is 0.017 and 0.015 (determined from gaussian fitting),
respectively. From the residual time series shown in Fig. 1e in Namekata et
al. (2021) the depths are in the order of 0.015. The maximum projected
velocities of the absorption reaches $\sim$-600 .. $\sim$500 km s-1 and
$\sim$-500 .. $\sim$500 km s-1 for our first two residual spectra, the
corresponding residual spectra in Fig. 1e from Namekata et al. (2021) have
projected velocity ranges of $\sim$-500 .. $\sim$300 km s-1 and $\sim$-200 ..
$\sim$400 km s-1. Our velocities are determined from Gaussian fitting, the
ones from Namekata et al. (2021) by eye from the non-fitted residual spectra.
Considering the fact that our temporal resolution and S/N is significantly
lower than the observations presented in Namekata et al. (2021) the depth and
also the projected velocity ranges agree well. Our observations reveal the
second half of the evolution of the eruptive filament structure until the
signature vanishes and therefore our observations complement the observations
from Namekata et al. (2021) and result in a full coverage of a highly
interesting stellar activity event.
Namekata et al. (2021) interpreted this event as probable filament eruption
following a flare. According to the high projected velocity of the bulk
material being close to the escape velocity of the star we support this
interpretation. From our observations one can see that the event lasted 1.5
hours longer than presented in Namekata et al. (2021), simply because their
spectral time series ended. Already in their observations the bulk velocity
shifted towards the red, meaning that hydrogen plasma is moving towards the
star. In our third residual spectrum (see Fig. 3) which includes already times
when the spectral time series presented in Namekata et al. (2021) had ended,
reveals that the projected velocity of the back-falling material stays at the
same level as seen in the last residual spectrum in Namekata et al. (2021).
The following residual spectrum in Fig.3 (no.4) shows then already a reduced
projected velocity until in residual spectrum no.5 the absorption has zero
velocity. In residual spectrum no.6 the absorption signature is then gone.
This behaviour is consistent with the morphology presented in Namekata et al.
(2021).
The other flare events shown in the present study (see Fig. 11 and Fig. 12) do
not show blue wing enhancements or absorptions, i.e. no accompanied
filament/prominence eruption is detected with those flares. The morphology of
the flares follow a typical flare light curve with a faster rise and a slower
decay, except for one event, namely event 04/04/2020. This one reveals no
typical H$\alpha$ core enhancements possibly caused by H$\alpha$ footpoint
brightenings due to fast electron

Figure 9: Upper panel: TESS light curve (black solid line) and H$\alpha$ light
curve (blue dotted line) extracted from the OLG data for event 19/02/2020.
Middle panel: TESS light curve and H$\alpha$ light curve extracted from the
OLG data for event 04/04/2020. Lower panel: TESS light curve and H$\alpha$
light curve extracted from the OLG data for 14/04/2020.

beams heating the chromosphere. We only see variation in two spectra, those
are spectra no.14 and 15, which show red shifts of 100 and 200 km s-1,
respectively. In the following spectrum no extra emission is visible anymore.
These red-shifted extra emissions could be caused by coronal rain in cool
H$\alpha$ flare loops (Wollmann et al., 2023), especially as these asymmetries
are visible in the decaying tail of the flare.
A distinct and long-lived (for five residual spectra, i.e. for 2.5 hr) red
asymmetry is also visible in event 22/04/2020 (see right panels of Fig. 12).
In residual spectrum no. 13 a broad symmetrical emission is visible which
slowly shifts to the red with a maximum bulk velocity of 150 km s-1 detected
in residual spectrum 17. In residual spectrum no. 18, which is the last of the
series, still an enhancement is visible but the residual spectrum is also
rather noisy, which does not allow a distinct identification of an
enhancement.
The other two flares, events 04/04/2020 and 14/04/2020 (see Fig. 11), reveal
only candidates of minor red asymmetries. We use the term candidate here
because the red asymmetries are very subtle and similar subtle emissions are
seen as fluctuations in the residual spectra of the series. Event 04/04/2020
is visible in two residual spectra, i.e. 1 hr, whereas event 14/04/2020 is
visible in 8 residual spectra which corresponds to 4 hr. The spectral residual
peaks are partly shifted by one wavelength bin to the red for the flare events
04/04/2020 and 14/04/2020. The active spectra reveal a core enhancement
together with enhanced H$\alpha$ red wings relative to the symmetric quiescent
H$\alpha$ profile. There seems to be a plasma component causing the red
neighbouring wavelength bin to be somewhat larger than the H$\alpha$ core
wavelength bin relative to the quiescent spectrum. Similar red asymmetric
spectral residual peaks have been shown only recently in dMe star flare
residual spectra by Notsu et al. (2024), who performed an extensive analysis
of a number of dMe star flares spectroscopically. Such red asymmetries have
been also modelled by Wollmann et al. (2023) and have been attributed to
downward flowing material in flaring loops termed coronal rain.
Only very recently Namekata et al. (2024) present again spectroscopic
H$\alpha$ observations of EK Dra, this time with partially coordinated X-ray
observations as well as broadband photometry from TESS. The authors find three
superflares from which two were accompanied with eruptive prominences, the
projected bulk velocities are fast (few hundreds of km s-1) but significantly
below the escape velocity of EK Dra. Together with their previous studies
(Namekata et al., 2021, 2022a) this results in total to five superflares from
which three were accompanied by filament/prominence eruptions. These projected
bulk velocities are lower limits, so there is a chance that these
filaments/prominences have indeed left the star. The superflares show
broadband (TESS) energies ranging from 2 … 26$\times$1033 erg and H$\alpha$
energies ranging from 2 … 49$\times$1031 erg. The authors give an association
rate of 60% between superflares and eruptive filaments/prominences, so three
out of five events show possible accompanying eruptive prominences/filaments.
We have captured four TESS superflares spectroscopically in H$\alpha$ and one
late evolution of a possible filament eruption, so an association rate of 20%
or one out of five events (although we have not observed the superflare
related to the possible filament eruption, but we know from TESS that there
was one). Our superflares show broadband energies in the range of
2-7$\times$1033 erg and H$\alpha$ energies in the range of 7-80$\times$1031
erg. We have one superflare (event 14/04/2020) in the sample which shows twice
the energy of the most energetic superflare presented in Namekata et al.
(2024), mainly because the duration of the flare lasts for four hours, whereas
their most energetic H$\alpha$ flare has a duration of about two hours. In
Namekata et al. (2024) the Seimei telescope is used, a 3.8 m telescope which
allows observing at a much higher cadence (1-2 min) than we are able with our
0.5m aperture. Therefore we may miss short-lived ($<$30 min) flares and
filaments/prominences. The spectral resolving power of the spectrographs is
similar. Only one superflare of the sample presented in Namekata et al. (2024)
shows no asymmetries, three show blue asymmetries and one shows red
asymmetries. The superflares presented in the present study show red
asymmetries only.
We have captured the time evolution of three TESS flares completely. These are
events 19/02/2020, 04/04/2020, and 14/04/2020. For these three flares we
overplot the TESS light curves with the OLG H$\alpha$ light curves extracted
from the optical spectra. In Fig. 9 one can see that the flare evolution in
both wavelength bands, narrow and broad, is different. The H$\alpha$ light
curves suffer from a low cadence, and therefore there is some uncertainty in
the peak time of the H$\alpha$ light curves, but the H$\alpha$ flare decay is
significantly different from the TESS broadband light curves. The H$\alpha$
decaying flare tails are more extended in time than their TESS counterparts.
For event 19/02/2020 and event 14/04/2020 this is likely related to the
Neupert effect which is known from solar flares (Neupert, 1968; Hudson, 1991).
White light flares are likely caused by fast electron beams accelerated from
reconnection regions in solar flares, which hit denser solar atmospheric
layers. H$\alpha$ flares are likely caused by electron beams which dissipate
their energy in the chromosphere. H$\alpha$ flare emission in solar flares is
also a product of post-flare loops occurring in the gradual flare phase. which
radiate also in H$\alpha$. In the impulsive flare phase chromospheric
evaporation (upward motion of plasma within the flare loops, e.g. Heinzel et
al., 1994) can occur and also chromospheric condensation can occur (downward
motion, e.g. Ichimoto & Kurokawa, 1984; Namekata et al., 2022b). For white
light flares it has been also suggested that not only foot point emission may
cause the main emission but that also, especially on young main-sequence
stars, the large flaring loops may contribute significantly to white light
emission in stellar flares (see Heinzel & Shibata, 2018). But there are also
flares on stars where white light and H$\alpha$ share a similar evolution.
E.g. Namekata et al. (2021) present a TESS white light flare which shows
similar temporal evolution as the H$\alpha$ flare. The same applies to the
flare which was accompanied by the filament eruption on EK Dra presented in
this study. As Namekata et al. (2021) captured also the flare
spectroscopically, they found that the white light and H$\alpha$ evolution of
the flare are similar.
Event 04/04/2020 differs in the H$\alpha$ light curve from the other two. This
event shows no typical H$\alpha$ light curve. Actually we see only two light
curve bins of enhanced flux in the middle panel of Fig. 9. These occur in the
decaying tail of the TESS flare in a form of central weak excess emission
(residual spectrum no.14 in the left panel of Fig. 12), followed by a
distinctively red-shifted excess emission. The first broad and central extra
emission may be attributed to H$\alpha$ footpoint emission whereas the
following red-shifted and stronger excess emission may be attributed to plasma
in post-flare loops. As young main-sequence stars may have much larger flare
loops as on the Sun, the loop excess emission may be possibly larger than the
excess emission originating from the footpoints. Mullan et al. (2006) deduced
lengths of flare loops based on EUV/X-ray observations of a number of stars
including EK Dra. These authors analysed two flares on EK Dra and found loop
lengths of 0.28 and 0.42 L/R⋆ (ratio of loop length to stellar radius),
respectively, which is in the range of larger flare loops on the Sun (see e.g.
Reale et al., 1997; Shibata & Magara, 2011).
We found spectroscopic signatures of flares and a filament eruption on EK Dra,
but why did we find no activity signatures on the other sample stars? The
answer lies already in the TESS light curves. One TESS sector has a length of
$\sim$ 27 days on average. The total spectroscopic observing times per star
are $\sim$14 days for EK Dra, $\sim$11 days for $\pi^{1}$ UMa, $\sim$14 days
for $\chi^{1}$ Ori, $\sim$28 days for HN Peg, and $\sim$1 day for $\kappa^{1}$
Cet. EK Dra was so far observed in ten TESS sectors (14, 15, 16, 21, 22, 23,
41, 48, 49, 50) and has revealed flares in every TESS sector; $\pi^{1}$ UMa
was observed so far in three TESS sectors (20, 47, 60) and has revealed so far
only one flare, $\chi^{1}$ Ori was observed in so far three TESS sectors (43,
44, 45) and has revealed three flares in one TESS sectors, in the remaining
two sectors no flares were detected; HN Peg was so far observed in one TESS
sector (55) and revealed no flares; $\kappa^{1}$ Cet was observed so far in
two TESS sectors (4,31) and revealed no flares (see Fig. 1 for TESS light
curves of the target sample stars). For this comparison between total
observing time and number of TESS flares, we set white light flare equal to
H$\alpha$ flare, which is not totally correct but a reasonable assumption on
the basis of the results presented in Kretzschmar (2011); Watanabe et al.
(2017), where a significant fraction ($>$50%) of the flares show white light
components based on solar data sets. If we multiply the TESS flare rate of our
stars (see Table 1) with their H$\alpha$ observing time then we get the number
of flares which should have been detected. For EK Dra this gives 6, for HN Peg
$<$3, and for the remaining stars 0 flares, which is roughly consistent to
what we detected.
In Table 3 we list the expected (see also Table 1) H$\alpha$ flare rates
exceeding an energy of 1032 erg which is a reasonable energy cut-off according
to the estimated H$\alpha$ flare energies given in Table 2. Only for EK Dra we
can compare expected and observed flare rate and here we see that the expected
rate is about an order of magnitude higher than the observed one. Also if we
try to relate the observed upper limits of the remaining stars with the
expected flare rates is of limited usefulness, as the upper limits are always
strongly dependent on the total observing time invested. One can of course
derive observed flare and CME rates from very few detected events, as done in
this study, but we have to caution that those are of low significance, as
those are determined from very few events.

Table 3: H$\alpha$ flare rates of the present study (3rd column) in comparison to the expected H$\alpha$ flare rates (2nd column). Only EK Dra has an observed H$\alpha$ flare rate, for the remaining stars we can give only upper limits. name | H$\alpha$ flare rate | H$\alpha$ flare rate
---|---|---
| expected | upper limit/observed
| N(>1032erg) day-1 | [day-1]
EK Dra | 2.81 | 0.29${}^{0.229}_{-0.139}$
HN Peg | 0.48 |  $<$0.11
$\pi^{1}$ UMa | 0.46 |  $<$0.21
$\chi^{1}$ Ori | 0.36 |  $<$0.28
$\kappa^{1}$ Cet | 0.23 |  $<$2.22

Another point to consider in the discussion of non detections is that for
ground-based observations one will not get uninterrupted data sets such as for
satellite observations, due to obvious reasons. This makes then the detection
of sporadic events a game of chance. We missed all stronger flares in the EK
Dra TESS light curves (see Fig. 10), but we still caught some flares because
EK Dra shows several flares in one TESS sector, but e.g. for $\chi^{1}$ Ori or
$\pi^{1}$ UMa, which show one or two flares in one sector, then in other
sectors no flares, the detection is challenging. So summarizing, especially
based on the TESS light curves, its seems reasonable that we detected
signatures of activity on EK Dra and no signatures on HN Peg although we
observed this star twice as much as EK Dra.
Solar analysis: The solar data analysis was used to study the detectability of
flares and filament eruptions on stars being part of the observed sample,
because we did not detect any signature of flares and CMEs on all of the
target stars besides EK Dra. Namekata et al. (2021) used also solar data, in
their study from Hida observatory, to base their interpretation of a very
probable stellar filament eruption on morphologies of solar filament
eruptions. Otsu et al. (2022) investigated a handful of solar activity
phenomena (flares, eruptive filaments/surges, and prominences) of the Sun seen
as a star, to better interpret stellar observations. Here, we use a complex
solar activity event involving flares, filament eruptions, and back flowing
material, not to compare its morphology but to evaluate its detectability in
OLG spectra.
The parameters of the solar event used in the present study are representative
for solar H$\alpha$ flares when compared to statistical studies of H$\alpha$
flares on the Sun (see e.g. Temmer et al., 2001). The H$\alpha$ pendant of the
X-class flare (see Fig. 5) has a rise time of $\sim$3 min and a duration of
$\sim$20-40 min, this agrees well with typical solar H$\alpha$ flare
parameters. When comparing this event with the event which occurred on EK Dra
(see Fig. 3), we see that the process of backflowing material on the Sun lasts
for $\sim$ 40 min whereas on EK Dra we identify two spectral residuals (no.3,
4) which reveal bulk absorptions with red velocities and therefore a duration
of one hour, which is in agreement with the solar event.
If we consider stellar integration times (3 min, 30 min), then the solar Sun-
as-a-star signature changes its appearance dramatically. This is shown in the
lower panel of Fig. 8 where the same event is shown as in the upper panel,
namely the EWs of the solar integrated signatures superimposed on a typical
spectrum of HN Peg obtained at OLG, but with integration times an order of
magnitude larger, i.e. around half an hour, which is the same integration time
which was needed to make the filament eruption visible on EK Dra. What one can
see clearly is that due to the much lower temporal resolution the signatures
falling in one time bin add up and are not well identifiable anymore. The
first time bin represents the pre-flare phase. The second time bin includes
pre-flare and impulsive flare phase and also a part of the decaying tail. The
third time bin then includes the second flare, occurring in the decaying tail
of the first flare, and also the major filament eruption including back
flowing material. The fourth time bin includes the post-flare phase which
interferes already with the next pre-flare phase. The second time bin includes
also the first filament eruption which remains completely invisible, but also
in the three minute averages (see upper panel of Fig. 8), this eruption is not
visible as it interferes with the impulsive phase of the flare.
We have constructed the solar signature from a fixed number of solar pixels
covering the active region from where the activity signatures originate. To
enhance the strength of the individual signatures one may have isolated those,
preventing thereby that the simultaneously occurring activity signatures (e.g.
flare and filament signature) may interfere and thereby affect the observed
signature. This is certainly worth to be investigated, as of course flares can
occur also without any eruptions. However as we investigated the event on EK
Dra which revealed a flare, a filament eruption and back-flowing material, we
were interested in a similar solar event.
To evaluate the detectability of activity signatures we have decided to use
two approaches. In approach 1 we utilize the same methodology as used for the
stellar data, namely to identify signatures of flares and/or erupting
filaments/prominences in residual spectra (see Fig. 6). In approach 2 we
determine the EW from the spectra and identify signatures of flares and/or
filaments/prominences in the EW time series (see Fig. 8). Approach 2 only
allows to distinguish between emission and absorption features. With the EW
one is not able to distinguish between blue or red absorption, i.e. erupting
or back-flowing material. For an estimation on detectability this is
sufficient. For the investigation of direction dependent flows one has still
to investigate the spectra themselves which we have shown with approach 2. To
make the signatures visible in Sun-as-a-star observations we had to scale the
solar observations in area. In general, the less pronounced the feature the
larger the area ratio scaling needs to be to make the signature visible. Using
approach 1 the flare is visible with scaling any area and the filament
eruption becomes visible already with an area scaling of 18%. Using approach
2, already with an area ratio of 12% the flare peak becomes visible and with
an area ratio of 24% the filament eruption/back falling material becomes
visible. Approach one is more sensitive to area ratio scaling than approach 2
and therefore also more useful to be used to search for activity signatures on
stars.

##  5 Summary and Conclusions

We have presented optical spectroscopic monitoring from OLG of a small sample
of young solar analogues. The observations revealed a very low level of
detectable activity in H$\alpha$ for the majority of stars in the sample
except for EK Dra, the youngest star in the sample. In the period from January
to April 2020 we have detected spectroscopically four flares and one episode
of a filament eruption on EK Dra. The eruptive event has been also published
by Namekata et al. (2021) who observed EK Dra at the same time. We have
started observing when the eruptive event was already ongoing and captured the
late stages of the eruption and the back falling material until the signature
vanished. With our observations from OLG we could complement the erupting
filament revealing its full evolution. In the typical exposure times used for
spectroscopic monitoring at OLG (three minutes) we did not see the absorption
signature caused by an eruptive filament on EK Dra. Building half hour
averages revealed the absorption signatures. To estimate the detectability of
solar activity events on our target stars, we spatially integrate over an
active solar region revealing a complex event including flares, filament
eruptions and back falling material observed by MCCD on MSO, similar to the
event on EK Dra. As expected, when superimposing the spectral spatially
integrated solar residuals of a solar event on a typical OLG spectrum of e.g.
HN Peg then it remains invisible. One needs to increase the event area to make
the signature visible in a spectrum. We need to scale the active region with
factors resulting in a fractional area of the active region being 18% in
residual spectra and 24% in equivalent width time series, respectively, to
make the signature detectable above the noise. This behaviour is consistent
with the fact that younger stars have larger active areas and therefore those
can be detected in stellar spectra, whereas solar active regions are too small
relative to the solar disk to be seen in full disk integrated light i.e. Sun-
as-a-star observations. However, even on the other stars of the sample,
activity signatures had been expected (from the H$\alpha$ flare rates) but
were not visible in the data. We therefore conclude that on solar-like stars
already in the first few hundred of Myr the occurrence rates of more massive
eruptive filaments/prominences decreases significantly. With our observational
setup we might have detected massive events only.
The intention of this study was the statistical determination of parameters of
stellar eruptive filaments/prominences and their relation to flares. We found
four flares and one filament eruption on one star. We know that the filament
eruption was accompanied by a flare but the other four flares did not show
signatures of filament eruptions. So one out of five flares on EK Dra shows an
accompanying filament eruption, but this result is far from being
statistically significant. Although the observational efforts have been
increased in the past few years to detect stellar CMEs still the number of
distinct events is low. We know many more candidate events, at least for the
method of Doppler-shifted absorption/emission (e.g. Fuhrmeister et al., 2018;
Vida et al., 2019). One way to obtain statistics is to focus on the numerous
(few hundreds) candidate events and try to better understand those. This has
already partly begun with the systematic investigation of spatially integrated
solar, i.e. Sun-as-a-star signatures of flares and eruptive filaments, with
the aim to better characterize stellar signatures of flares and eruptive
filaments/prominences, including their temporal evolution (see e.g. Leitzinger
et al., 2021; Namekata et al., 2021; Leitzinger, 2022; Otsu et al., 2022).
These studies used solar instruments capable of spatially resolved 2D
spectroscopy (such as MCCD on Mees Solar Observatory/MSO) or full-disk
photometry in various filters (Solar Dynamics Doppler Imager/SDDI on The Solar
Magnetic Activity Research Telescope/SMART).
With this study we have demonstrated that small-sized telescopes can be used
to infer spectroscopic activity signatures on bright solar-like stars. From
the spectroscopic monitoring presented in this study we have seen that CMEs,
more energetic and massive than occurring on our present-day Sun, on few
hundred Myr old solar analogues are not a frequent phenomenon.

## Acknowledgements

We thank the anonymous referee for very valuable comments which improved the
manuscript. This research was funded in whole, or in part, by the Austrian
Science Fund (FWF) [10.55776/P30949, 10.55776/I5711]. For the purpose of open
access, the author has applied a CC BY public copyright licence to any Author
Accepted Manuscript version arising from this submission. This paper includes
data collected by the TESS mission. Funding for the TESS mission is provided
by the NASA’s Science Mission Directorate. Mees acknowledgements go to NASA
HDEE Grants 80NSSC18K0064 and 80NSSC18K1658 for the data rescue effort, NASA
Grant NAGW 1542 for the instrument fabrication. Support was also provided from
Lockheed under NASA contract NAS8-37334 with Marshall Space Flight Center and
the Yohkoh mission contract NAS8-40801. This work presents results from the
European Space Agency (ESA) space mission Gaia. Gaia data are being processed
by the Gaia Data Processing and Analysis Consortium (DPAC). Funding for the
DPAC is provided by national institutions, in particular the institutions
participating in the Gaia MultiLateral Agreement (MLA). The Gaia mission
website is https://www.cosmos.esa.int/gaia. The Gaia archive website is
https://archives.esac.esa.int/gaia.

## Data Availability

For the present study we used data from TESS, Mees, GAIA DR3, and OLG. TESS
data are freely available via the Mikulski Archive for Space Telescopes (MAST,
https://mast.stsci.edu/portal/Mashup/Clients/Mast/Portal.html). Mees data are
freely avalailable and accessible under https://www.cora.nwra.com/MCCD/. Gaia
data are freely available and accessible under
https://archives.esac.esa.int/gaia. OLG data are not freely available but can
be made available under request via email.

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##  6 Appendix

Figure 10: TESS light curves of EK Dra during which OLG spectroscopic
observations are available. As one can see we missed all larger flares. Only
few weaker flares (although all of them are superflares) have been captured.
The numbers given indicate the spectroscopic data sets per month for the
detected events so that the reader can relate the TESS light curve with the
residual spectra in Fig. 3, 9, 11, 12.

Figure 11: Photometric (upper panel) and residual spectral time series (lower
panel; from top to bottom) of EK Dra from the night of the 4th of April 2020
(left panels) and of EK Dra from the night of the 14th of April 2020. Grey
shaded areas denote the 1-$\sigma$ uncertainty of the residuals. Red solid
lines are gaussian fits (single and double) to the excess emission visible
from residuals 14-15 (lower left panel) and residuals 2-9 (lower right panel).
Solid blue and green lines are the single components of double gaussian
fitting.

Figure 12: Photometric (upper panel) and residual spectral time series (lower
panel; from top to bottom) of EK Dra from the night of the 19th of February
2020 (left panels) and of EK Dra from the night of the 22nd of April 2020.
Grey shaded areas denote the 1-$\sigma$ uncertainty of the residuals. Red
solid lines are gaussian fits (single and double) to the excess emission
visible from residuals 3-6 (lower left panel) and residuals 13-17 (lower right
panel). Solid blue and green lines are the single components of double
gaussian fitting.