Probable detection of an eruptive filament from a superflare on a solar-type star

Solar flares are often accompanied by filament/prominence eruptions (~104 K and ~1010−11 cm−3), sometimes leading to coronal mass ejections that directly affect the Earth’s environment1,2. ‘Superflares’ are found on some active solar-type (G-type main-sequence) stars3–5, but the filament eruption–coronal mass ejection association has not been established. Here we show that our optical spectroscopic observation of the young solar-type star EK Draconis reveals evidence for a stellar filament eruption associated with a superflare. This superflare emitted a radiated energy of 2.0 × 1033 erg, and a blueshifted hydrogen absorption component with a high velocity of −510 km s−1 was observed shortly afterwards. The temporal changes in the spectra strongly resemble those of solar filament eruptions. Comparing this eruption with solar filament eruptions in terms of the length scale and velocity strongly suggests that a stellar coronal mass ejection occurred. The erupted filament mass of 1.1 × 1018 g is ten times larger than those of the largest solar coronal mass ejections. The massive filament eruption and an associated coronal mass ejection provide the opportunity to evaluate how they affect the environment of young exoplanets/the young Earth6 and stellar mass/angular momentum evolution7. An energetic eruptive filament on EK Draconis most probably launched a coronal mass ejection with a mass ten times larger than the largest solar coronal mass ejection. Studying such ejections provides insight into stellar angular momentum loss and the habitability of orbiting planets.


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Nature astroNomy (Fig. 1c-e and Extended Data Figs. 2a and 3a). Both ground-based spectroscopic observations simultaneously recorded the same spectral change, demonstrating that low-temperature and high-density neutral plasma above the stellar disk moves at high speed toward the observer before some parts finally start to fall back to the surface. In addition, the deceleration is not monotonic: it was 0.34 ± 0.04 km s −2 in the initial phase, dropping to 0.016 ± 0.008 km s −2 in the later phase (Fig. 1c,d and Extended Data Fig. 3b). This is interpreted in terms of changes in the height of the ejected mass. The observed deceleration is in good agreement with that due to the surface gravity of approximately 0.30 ± 0.05 km s −2 (ref. 9 ), although the initial value is slightly larger.
How much do the stellar spectral changes obtained here actually resemble those of solar filament eruptions? Blueshifted Hα absorption profiles are often observed from solar filament eruptions 1,14 . As in Fig. 2, we generated spatially integrated Hα spectra of a solar flare/filament eruption that occurred on the solar disk using the SMART (Solar Magnetic Activity Research Telescope) data 15 (Extended Data Fig. 4 and Supplementary Video 1). We converted to the full-disk pre-flare-subtracted spectra by multiplying by the partial-region/full-disk ratio (that is, virtual Sun-as-a-star spectra). We found that the blueshifted absorption component at approximately 100 km s −1 was predominant soon after the solar flare, and the spatially integrated Hα EW showed enhanced absorption (Fig. 2a). These blueshifted profiles are unequivocally due to the filament eruption. Later, the blueshifted component decelerated and gradually turned into slow, redshifted absorption (Fig. 2b,c). The Hα EW returned to the pre-flare level in approximately 40 min (Fig. 2a). Although the energy scales and velocities are different, the solar data strongly resemble the spectral changes in the superflare on EK Dra (see Supplementary Information for another event). This similarity suggests that the stellar phenomenon is the same as the simply magnified picture of the solar filament eruption.
A filament eruption is the only explanation for the blueshifted absorption component on EK Dra by solar analogy 1 . The hypothesis that the blueshifted absorption on EK Dra might come from up-/downflow in flare kernels must be rejected because they never show Hα absorption 16,17 . Also, downflow in cooled magnetic loops (known as post-flare loops) 14 shows redshifted absorption, so they cannot explain the blueshifted absorption. (However, the redshifted absorption in EK Dra in the later phase might be caused by post-flare loops 14 .) Rotational visibilities of prominences or spots also are not adequate to explain it, since the rotation speed of EK Dra is only 16.4 ± 0.1 km s −1 (ref. 9 ). Thus, we concluded that we detected a stellar filament eruption on the solar-type star.
Some observational signatures for stellar filament eruptions or CMEs have been reported previously for cooler K-M dwarfs [18][19][20][21][22] and evolved giant stars 23 (see Methods and refs. 6,24 for reviews). Velocity (km s -1 ) Velocity (km s -1 ) Velocity (km s -1 ) 1,000 The 1σ value of the pre-flare light curve (−150 min to 0 min) is shown in blue. b, Light curves of the Hα EW observed using the medium-dispersion spectrograph MALLS (Medium and Low-Dispersion Long-Slit Spectrograph) at the Nayuta telescope (grey circles) and the low-dispersion spectrograph KOOLS-IFU (Kyoto Okayama Optical Low-Dispersion Spectrograph with optical-fibre Integral Field Unit) installed at the Seimei telescope (red triangles) during the same observing period as in a. The Hα emissions were integrated within ±10 Å from the Hα line centre (6,562.8 Å) after dividing by the continuum level, and the pre-flare level was subtracted. The positive and negative values represent emission and absorption, respectively, compared with the pre-flare level. The 1σ value of the pre-flare light curve (−150 min to 0 min) is plotted in red and black for Seimei and Nayuta data, respectively. c,d, Two-dimensional Hα spectra obtained using the Seimei telescope (c) and the Nayuta telescope (d). The red and blue colours correspond to emission and absorption, respectively. The dashed lines indicate the stellar surface gravity (g * ) and half of the surface gravity (0.5 g * ). c and d share the upper colour bar. e, Temporal evolution of the pre-flare-subtracted Hα spectra observed using the Seimei telescope (red) and the Nayuta telescope (black), with the spectra shifted by constant values for clarity. The spectra are binned in time, and the integration periods correspond to the horizontal axes of a-d. The intensities are normalized by the stellar continuum level. The vertical dotted line indicates the Hα line centre, and the horizontal dotted lines indicate the zero levels for each spectrum. The 1σ error bar around the line core, based on the residual scattering in the line wing, is also shown.

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The observation of a giant star shows a blueshifted X-ray emission line of 90 km s −1 in the post-flare phase and hotter CME is proposed as a possible explanation 23 . Recently, X-ray/extreme UV dimmings have been reported as indirect evidence of stellar CMEs on K-M dwarfs 22 . In M-dwarf flares, many blueshifted Balmer/UV line emission components have been reported [18][19][20][21]24 , which are interpreted as filament eruptions. Some M-dwarf flares share properties with the eruption on EK Dra: the blueshift emissions have high velocities of hundreds of kilometres per second, and some exhibit velocity changes and appear after the impulsive phase 20,21 . For M-dwarf events, the number of studies reporting highly time-resolved velocity variations of blueshift components is still insignificant (~5 min cadence), and a simultaneous white-light flare has never been detected. Our detection of a stellar filament eruption is reliable because we provided solar counterparts, highly time-resolved spectra (~50 s cadence) and a simultaneous TESS white-light flare. What properties does the filament eruption on EK Dra have? The maximum observed velocity of the blueshifted component was ~−510 km s −1 with a width of 220 km s −1 . This is larger than the typical velocities of solar filament eruptions (10-400 km s −1 ) associated with CMEs 2 , although it is a little smaller than the escape velocity at the surface on EK Dra (~670 km s −1 ). The cool plasma reached at least ~1.0 stellar radius from the stellar surface (or the initial height) as derived by integrating the velocity over time (or ~3.2 stellar radii from the stellar surface on the basis of the deceleration rates). In this case, a projection angle of at most 45° can be allowed when we assume that the event occurs on the disk centre. On this projection angle, the velocity can be up to ~−720 km s −1 , so there is a possibility that the velocities of some components of the EK Dra eruption could exceed the escape velocity. However, it should be noted that there are weak redshifted components with a velocity of a few 10 km s −1 in the late phase, indicating that some materials fell back to the star. This is often observed in the case of solar filament eruptions with CMEs 25 .
The filament area is estimated to be 1.6 × 10 21 cm 2 (5.6% of the stellar disk), and the erupted mass is calculated to be 1.1 +4.2 −0.9 × 10 18 g on the basis of the absorption components. The mass is more than ten times larger than those of the largest solar CMEs 26,27 (it should be noted that the mass can be somewhat under-/overestimated; Methods). This mass estimate is in reasonable agreement with those predicted from empirical 26,27 and theoretical 28 solar scaling relations between CME mass and flare energy within the error bars (~9.4 +3.2 −2.4 × 10 16 and 3.1 +1.6 −1.1 × 10 17 g for refs. 27 and 26 , respectively) ( Fig. 3a). This suggests that the stellar filament eruption can share a common underlying mechanism with smaller-scale filament eruptions/CMEs (that is, magnetic energy release 1,28 ) although the absolute values of most physical quantities are very different.
Moreover, the kinetic energy is calculated to be 3.5 +14.0 −3.0 × 10 32 erg, which is 16% of the radiation energy in white light. The magnetic energy stored around the starspots on EK Dra can be at least 8.0 × 10 35 erg, which is enough to produce superflares and filament eruptions with energy of ~10 33 erg. In addition, this value is slightly smaller than those extrapolated from the solar CME scaling law (4.8 +1.1 −0.9 × 10 33 erg; ref. 27 ) (Fig. 3b), which is similar to the filament eruption/CME candidates on other stars 24 . In previous studies, it has been argued that kinetic energy can be reduced by overlying magnetic fields 24,29 . The deceleration of our events was a few times 10% larger than the stellar gravity (Extended Data Fig. 3b). The strong magnetic fields on EK Dra have been reported before 9 and may support the above explanations. However, its small kinetic energy can also be understood through a solar analogy: the velocities of (lower-lying) filament eruptions are usually four to eight times lower than those of the corresponding (higher-lying) CMEs 2 , and therefore the kinetic energies of filament eruptions are typically smaller (green symbols in Fig. 3b).
Did a CME occur in this event? Obviously, the line-of-sight velocity of ~510 km s −1 was lower than the escape velocity and some masses fell back, which may indicate a so-called 'failed' filament eruption 29 . However, this does not necessarily mean that a CME did not occur, again by solar analogy. In fact, the erupted filaments often fall back to the Sun when CMEs happen. For example, a well studied solar event on 7 June 2011 involved a 200-600 km s −1 filament eruption where much filamentary material fell back to the Sun, but some mass clearly escaped as a CME with velocities of ~1,000 km s −1 (ref. 25 and Supplementary Information). The event

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Nature astroNomy on EK Dra may correspond to this solar event. In addition, ref. 30 showed that whether a solar filament eruption leads to a CME can be simply distinguished by a parameter of (V r_max /100 km s −1 ) (L/100 Mm) 0.96 , where V r_max is the maximum radial velocity and L is the length scale (Fig. 4). When the parameter is more than ~0.8, the probability that a filament eruption leads to a CME is more than 90% (ref. 30 ). The value of the parameter of eruption on EK Dra is ~18, meaning that our detection of the fast and sizable stellar filament eruption is indirect evidence that mass escapes into interplanetary space as a CME.
Finally, we summarize future directions of our findings (see Supplementary Information for details): It is speculated that the filament eruptions/CMEs associated with superflares can severely affect planetary atmospheres 6 . Our findings can therefore provide a proxy for the possible enormous filament eruptions on young solar-type stars and the Sun, which would enable us to evaluate the effects on the ancient, young Solar System planets and the Earth, respectively. Further, it is also speculated that stellar mass loss due to filament eruptions/CMEs can affect the evolutionary theory of stellar mass, angular momentum and luminosity 7,26 more importantly than can stellar winds. At present, frequency and statistical properties of CMEs on solar-type stars are unknown, but important insights into these factors will be obtained by increasing the number of samples in the future.  The red square represents the superflare on EK Dra, the black crosses denote solar CME data, the green triangles signify data for solar prominence/filament eruptions and surges taken from previous studies and the green plus sign is the solar filament eruption/surges displayed in Fig. 2 and Supplementary Fig. 9 (Velocity, mass, and kinetic energy: solar data), respectively ( . The kinetic energy of eruption on EK Dra is calculated to be 3.5 +14.0 −3.0 × 10 32 erg, which is outside the error range of the predicted value of 4.8 +1.1 −0.9 × 10 33 erg (ref. 27 ). The error bars are derived as the model errors (see 'Velocity, mass and kinetic energy: stellar data').

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Methods TESS light-curve analysis. TESS observed EK Dra (TIC 159613900) in its sectors 14-16 (18 July 2019-6 October 2020) and 21-23 (21 January 2020-15 April 2020). The TESS light curve from the 2 min time-cadence photometry was processed by the Science Processing Operations Center pipeline, a descendant of the Kepler mission pipeline based at the NASA Ames Research Center 12,31 . Extended Data Fig. 1 shows the light curve of EK Dra from BJD 2458945 (JD 2458944.997, 5 April 2020 11:56 ut; Sector 23), and the stellar superflare detected by TESS, the Seimei telescope and the Nayuta telescope in Fig. 1 is indicated with the red arrow in this figure. The quasiperiodic brightness variation is thought to be caused by the rotation of EK Dra with the asymmetrically spotted hemisphere 3,5 . The rotation period is reported as about 2.8 d (ref. 9 ). Although the superflare occurred near the local brightness maximum, some of the starspots are expected to be visible from the observer [32][33][34][35] . In Extended Data Fig. 1, other flares are also indicated using black arrows with more than two consecutive observational points whose flaring amplitude is more than three times the TESS photometric errors 3,36 . The white-light flare energy was calculated by assuming the 10,000 K blackbody spectra 36,37 (Flare energy). The pixel-level data analysis is shown in TESS pixel-level data analysis. The estimated occurrence frequency of superflares (>10 33 erg) in the TESS band was about once per 2 d, which means that about 12 nights' monitoring observations are necessary on average to detect one superflare from the ground-based telescope under a clear-sky ratio of 50%. This implies that our datasets are highly unique.

Spectroscopic data analysis.
Here, we present the utilization of low-resolution spectroscopic data from KOOLS-IFU 38 of the 3.8 m Seimei telescope 13 at Okayama Observatory of Kyoto University and MALLS 19,39 of the 2 m Nayuta telescope at Nishi-Harima Astronomical Observatory of the University of Hyogo. KOOLS-IFU is an optical spectrograph with a spectral resolution of R (λ/Δλ) ~ 2,000 covering a wavelength range from 5,800 to 8,000 Å; it is equipped with Ne gas emission lines for wavelength calibration and instrument characterization. The exposure time was set to be 30 s for this night. The sky spectrum was subtracted by using the sky fibres for each spectrum. The data reduction follows the prescription in ref. 40 . During this observation, the signal-to-noise ratio (S/N) for one frame is typically 172 ± 6. The observations using the Seimei telescope ended just after 133.7 min (Fig. 1b-d).
MALLS is an optical spectrograph with R ~ 10,000 at the Hα line covering a wavelength range from 6,350 to 6,800 Å; it is also equipped with Fe, Ne and Ar gas emission lines for wavelength calibration and instrument characterization. The sky spectrum was subtracted using a nearby region along the slit direction for each observation. The exposure time was set to be 3 min for this night. The MALLS data reduction follows the prescription in ref. 19 . The S/N for one frame is typically 86 ± 8 during this observation. For the MALLS data, the wavelength corrections are also performed for each spectrum using the Earth's atmospheric absorption lines.
We corrected the wavelength for the proper motion velocity of −20.7 km s −1 of EK Dra on the basis of Gaia Data Release 2 (ref. 41 ). Continuum levels are defined by a linear fit between the wavelength ranges of the Hα line wing (6,517.8-6,537.8 and 6,587.8-6,607.8 Å). We take the continuum level as the wavelength range between 6,517.8-6,537.8 and 6,587.8-6,607.8 Å to measure the EW (=∫(1 − F λ /F 0 ) dλ, where F 0 is the continuum intensity on either side of the absorption feature, while F λ represents the intensity across the entire wavelength range of interest). The original spectra are shown in Stability of pre-flare spectra. Extended Data Fig. 2 shows the pre-flare-subtracted Hα spectra during and after the superflare on EK Dra with higher time cadence than Fig. 1e. The narrowband Hα EW (Hα − 10 Å-Hα + 10 Å) is used for the measurements of the radiated energy and duration of the Hα flare because of the high S/N, and the broadband Hα EW (Hα − 20 Å-Hα + 10 Å) is used for the measurements of the amount of absorption (that is, mass and kinetic energy).

Solar data analysis.
In the main text, we showed the data of a C5.1-class solar flare (that is, the peak GOES soft X-ray flux F GOES is 5.1 × 10 −6 W m −2 ; hereafter 'Event 1' , see Table 1) and associated filament eruption around 07:56 ut, 7 July 2016, observed using the SDDI 15 Table 1). This paper used 70 min time series of the SDDI images taken from 07:30 ut on 7 July 2016 (Supplementary Video 1). As in Extended Data Fig. 4, the C5.1-class flare occurred around an active region, named 'NOAA 12561' , on the solar disk, and was accompanied by a typical filament eruption 15,42 . The spectra from the event are integrated over a spatial region that is large enough to cover the visible phenomena (the magenta region in Extended Data Fig. 4a,b). The spectra are reconstructed by using the template solar Hα spectrum convolved with the SDDI instrumental profile.
Here, we define L(λ, t, A) as the luminosity at a wavelength of λ and time of t that is integrated for the region A (that is, L(λ, t, A) = ∫ A I(t) dA; I(t) is intensity). We now define A local as the integration region (magenta region in Extended Data Fig. 4a,b), and A full-disk as the solar full disk. We first obtain the local (partial-image) pre-flare-subtracted spectra ΔS local , which are normalized by the local (partial-image) total continuum level (L(6,570.8 Å, t, A local )): where t 0 is a given time of the pre-flare period. Then, the (virtual) full-disk pre-flare-subtracted spectra ΔS full-disk are obtained by multiplying by the ratio of the partial-image continuum to full-disk continuum (total continuum ratio): and we obtain a virtual pre-flare-subtracted spectrum of this phenomenon as if we observed the Sun as a star. The EW of the Hα is also calculated using the ΔS full-disk , and we obtained the virtual Sun-as-a-star ΔHα EW (that is, differential Hα flux normalized by the full-disk continuum level).

Stellar velocity, mass and kinetic energy data.
For the stellar filament eruption, the velocity is derived by fitting the absorption spectra obtained using the Seimei telescope with the normal distribution N(λ, μ, σ 2 ), where μ is the mean wavelength and σ 2 is the variance. In Extended Data Fig. 3a, we plotted the temporal evolution of the velocity ((μ − λ)/λc, where λ is 6,562.8 Å and c is light speed) for the fitted absorption feature with the width of σ. We only plotted the data whose absorption we can expect that the filament is flying in our direction perpendicularly to some extent, so there would not be such a large difference between radial velocity and line-of-sight velocity. We expect that the radial velocity can be larger than the line-of-sight velocity if we assume the projection effect, while it will be about √2 times smaller at most if it erupts at a 45° tilt in the radial direction, which does not change our discussion. The solid line indicates the threshold that can roughly distinguish filament eruption with and without CMEs derived in ref. 30 . The threshold can be expressed as (V r_max /100 km s −1 )(L/100 Mm) 0.96 = 0.8, which is determined by using the Linear Support Vector Classification algorithm (see ref. 30 for the detailed method).

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features are clear enough to fit the shape with the threshold of the fitted absorption amplitude of >0.01 and fitted velocity dispersion of <500 km s −1 and >100 km s −1 . The threshold was determined by trial and error, and we find that many missed detections of absorption features occur when we select threshold values other than this one. The amplitude value of 0.01 corresponds to the detection limit when considering the typical S/N ~ 170 of the Seimei telescope/KOOLS-IFU, and the lower limit of 100 km s −1 is determined to avoid detecting the sharp noisy signals. About 27% of data points were discarded due to this threshold from the initial points (22 min) to final points (110 min), especially for the latter decaying phase. Here, the maximum observed velocity and its errors are calculated as 510 ± 120 km s −1 with a width of 220 ± 90 km s −1 from the mean values of the μ and σ of the first five points (t = 22-26 min in Fig. 1), respectively. The mean value of the velocity when the absorption becomes strong (t = 25-50 min in Fig. 1) is estimated as 258 km s −1 .
The plasma mass is simply calculated from the total Hα EW. We used the simple Becker cloud model 43 with optical depth at the line centre of the ejected plasma τ 0 of 5 (which is slightly more optically thick than solar filament eruptions; compare ref. 44 ), two-dimensional aspect ratio of 1 (that is, cubic), local plasma dispersion velocity W of 20 km s −1 and source function S of 0.1 on the basis of the solar observations 45 . The observed half-width of 220 km s −1 of the stellar blueshifted component is larger by one order of magnitude than the solar value, but here we use the solar value as a template. The dispersion velocity of 220 km s −1 is considered to be the upper limit of the local velocity dispersion because the ejected mass would have a complex two-dimensional velocity distribution, which can cause larger W in the integrated spectra. First, the modelled EW of enhanced absorption is calculated by using the Becker cloud model, when the plasma velocity v shift is −258 km s −1 , as where I 0λ is background intensity and I 0,Cont. is continuum intensity. This is the EW value for an extreme case when the full disk of the star is completely covered with absorbing, cool ejected plasma. By comparing the modelled EW (equation (3)) with the lowest observed stellar EW value of −0.16 Å (integrated for Hα − 20 Å-Hα + 10 Å; Supplementary Fig. 4c), the cool-plasma filling factor is calculated to be 5.9% of the stellar disk (that is, modelled EW/observed EW; area = 1.6 × 10 21 cm 2 ). Using the length scale of the ejected plasma, 3.9 × 10 10 cm (=area 0.5 ), the hydrogen column density is derived as 4.0 × 10 20 cm −2 from the assumed optical depth based on the plasma model 46 . In the model of ref. 46 , hydrogen/electron density is calculated by assuming an ionization equilibrium for a population of hydrogen atoms due to a balance between recombination and radiative photoionization through the Balmer/Lyman continuum. It should be noted that the ionization equilibrium of filaments on active stars may be somewhat different from the solar observations due to their high UV radiations, which may affect the evaluation of the mass of the ejecta. By multiplying the hydrogen column density by the filament area, we then obtained the plasma mass of 1.1 × 10 18 g. If the two-dimensional aspect ratio becomes 0.1, similar to a jet-like feature (x width:y width:z depth = 1:0.1:0.1), then the estimated mass becomes larger by a factor of 1.78. If optical depth ranges from 0.8 to 10 (ref. 44 ), the source function takes values of 0.02 or 0.5 and the dispersion velocity is 10 or 220 km s −1 (ref. 45 ), then the estimated masses change by a factor between 0.15 and 4.9. In Fig. 3a, we used the mass of 1.1 +4.2 −0.9 × 10 18 g for an optical depth of 5, and the uncertainties of the model (0. 15-4.9) are used as the error bars since the model-based errors are expected to be much larger than the observational errors. It should be noted that this mass estimate could be either a significant overestimate of the mass of an affiliated CME due to most of the filament falling back to the star, or a significant underestimate due to most of the CME actually being hot coronal material rather than cool filament. The plasma kinetic energy is then calculated as 3.5 +14.0 −3.0 × 10 32 erg by using the velocity of 258 km s −1 . The observed maximum velocity was 510 km s −1 in the early phase, so the kinetic energy can be larger by a factor of 4 although the absorption component was weak at that time.
A CME signature was reported from a blueshifted emission component of the cool X-ray O viii line (4 MK) in the late phase of a stellar flare on the evolved giant star HR 9024 23 . Although the time evolution of the blueshifted velocity is not obtained there, they detected the blueshifted emission component with a velocity of 90 km s −1 (escape velocity 220 km s −1 ) and interpreted it as a CME. The blueshifted plasma components at a few MK are also emitted from the upward flow in the confined flare loops (called 'chromospheric evaporation') in the case of solar flares, but they exclude the possibility considering that the other hotter lines do not show the blueshifted component in the post-flare phase. Although the spectral type of HD 9024 (evolved giant star) is very different from that of EK Dra and the velocity (90 km s −1 ) is smaller than our observation (510 km s −1 ), the two observations share the trend that the mass ejection signature is dominant in the post-flare phase.
Blueshifted emission components of chromospheric lines have been reported in association with Balmer-line flares mostly on active M/K dwarfs 18-21,48-57 (see refs. 24,44 for a summary). Time-varying blueshifted hydrogen emission components have also been reported with high time cadence on M dwarfs (for example, refs. 19,21 ). A similar case is reported for a UV flare on an M dwarf 20,70 . These may be evidence of stellar prominence eruptions/CMEs. It seems quite possible that the blueshifted emission lines on M dwarfs are closely analogous to the Hα absorption signatures studied in this Letter. The fundamental difference between G-dwarf and M-dwarf blueshift signatures is that for hotter G dwarfs Hα in an erupting filament will only be detectable in absorption, whereas for the cooler M dwarfs even the quiescent Hα line is in emission, so an erupting filament might be observed in emission as well (compare ref. 44 ). Blue-wing enhancements of M-dwarf flares are characterized by high velocity of several hundred kilometres per second (sometimes more than this) 18,53,55 , which cannot be explained by chromospheric evaporation flow associated with the chromospheric-line blueshift phenomenon observed in solar flares 16,44,[71][72][73][74] . The high velocities of M-dwarf flares are similar to that detected on EK Dra in this study (~510 km s −1 ). In addition, not all but some of the blueshift events on M dwarfs appear after the impulsive phase 20,21 , which shares properties with filament eruption events on EK Dra and the Sun in this study. Therefore, at present the blueshifted emission lines in M-type stars are most probably prominence eruptions.
In some cases of binary stars, eclipses of the white dwarf component have been interpreted as obscuration by stellar mass ejected from the late-type companion star 61,85 . Other than this, pre-flare dips have been reported in stellar flares, suggesting potential prominence eruptions/CMEs 86,87 . Radio observations have recently investigated the type II radio bursts associated with shocks in front of CMEs as possible indirect evidence of CMEs, but no significant signature has been obtained so far 47,62-65,67-69 . Recently, a stellar type IV burst event from the M-type star Proxima Centauri was reported and may be evidence for a stellar CME 56 .

Data availability
Source data are provided with this paper. In addition, all raw spectroscopic data are available either in the associated observatory archive (https://smoka.nao.ac.jp/ index.jsp for KOOLS-IFU data in Fig. 1 (available after January 2022); https://www. hida.kyoto-u.ac.jp/SMART/T1.html for some of the SDDI data in Fig. 2) or upon request from the corresponding author (for MALLS data in Fig. 1 and full raw data of SDDI). The TESS light curve is available at the MAST archive (https://mast.stsci. edu/portal/Mashup/Clients/Mast/Portal.html).