KOOLS-IFU ToO Observations toward
eROSITA-selected Tidal Disruption Event Candidates
Seimei Users Meeting, 12th August 2021

Taiki Kawamuro (Universidad Diego Portales/NAOJ)
A. Rau, I. Grotova, A. Malyali, A.Merloni, Z. Liu (MPE)
Y. Toba, S. Yamada (Kyoto U.), T. Nagao (Ehime U.)
Special thanks to H. Maehara (NAOJ) & K. Matsubayashi (Kyoto U.)

Tidal Disruption Events and X-ray Observations(NASA)(Komossa 2015)

Star

SMBH

• An invaluable class for studying the evolution of accreting matteraround SMBHs(Gezari et al. 2017)• Signature of a star being disrupted by a SMBH in galactic nuclei• Months-years lasting in X-ray and UV/optical as well
• Soft black-body X-ray emission, typical for ROSTA TDEs in 2000

kT = 42±0.7 eV
Γ = 3±1

eROSITA Survey for a New and Large X-ray TDE sample
SRG satellite
ART-XC• eROSITA:• Launch in December 2019
• ~ 0.5-2 keV (as main band)• ~ 10 times higher sensitivity than ROSAT• On-going all-sky scans (= eRASS)at a cadence of ~ 6 months-> 8 scans in total• 3 scans were already completed

eROSITA

Nuclear transients sample statistics(for one year b/w June 2020-June 2021)(Fig.: I. Grotova)

90

Flux distribution

80

Soft

70

50
40
30
20
10

1e-13

up)

• 61 soft candidates/yr, ~ 30 per eRASS

60

0

•
• after additional filtering (prior follow-

Soft421 Nuclear Transients

1e-13

1e-12
flux [0.2-5.0 keV]
1e-12

1e-11

1e-11

Nuclear transients sample statistics(for one year b/w June 2020-June 2021)(Gezari+21)• 421 Nuclear Transients• after additional filtering (prior followup)• 61 soft candidates/yr, ~ 30 per eRASS• Higher identification rates than recent
optical surveys !
(i.e., ~ 40 candidates/yr)

XMM/CXO archiveROSAT

eROSITA• Now is the time for systematic followup obs. toward X-ray TDEs
(Caution: this is until ~ 2024)

Brief Description of Search Procedure• Weekly search for new nuclear transients from galaxies without prior signsof AGN activity• Initial automatic selection based on• X-ray softness• Flare amplitude compared to previous eRASSs• Host galaxy properties or in pre-burst (e.g., W1-W2<0.5mag)• GAIA proper motion
X-ray/UV/Optical/Radio Follow-up• Verification and removal of contaminants• Redshift measurements• Light curve and spectral evolution• Energetics

kT~90eV

TDE

Hβ
Hγ

[OIII]FeXIV

AT 2019avd
cm-2 s-1 Å-1)

Hγ

HeII

FeXIV

Optical Spectroscopy Follow-up Eﬀorts
Massive eﬀort by many people combining a wide range facilities around the globe
LCO 2m + FLOYDS (MPE Time)• Queue-schedule access (~21 observations)

LCO 2mSALT

ANU 2.3m + WiFeS (G. Anderson, J. Miller-Jones, A. Kafka; AAL)
• Monthly observations (9 nights so far, ~30 targets)
SEIMEI

3.5m ESO NTT + EFOSC (PI: A. Malyali)• Guest observer program (48 targets)
3.8m SEIMEI (PI: T. Kawamuro, Y. Toba; IECs)• ToO Guest observer program (1 target)※ Only ToO access to spectroscopy
6.5m Magellan (PI: F. Di Mille, S. Ciroi; IECs)• Opportunity observations (12 sources)

ANU 2.3m

ESO NTT

Magellan~10m SALT (PI: D. Buckley; MoU)
• SALT Transient ToO project (4 sources)
Data stored on the SciServer

will enable us to discuss the accretion flow in terms of gas density, while confirming the decreasing trend.
● Evolution of Line Ratio:
Lastly, the ratio of He II 4686 and H↵ emission, often bright in TDEs (Figure 2), is another key tounderstand the accretion flow. Optically-selected TDEs sometimes showed weak Blamer emission lines inspite of bright He II emission, that is, the presence of strong UV emission ionizing H as well as He. vanVelzen et al. (2020) presented a trend that He II brighter TDEs would have more compact accretion flows,• Why
Seimei/KOOLS-IFU?
where
optically-thick
reprocessing
can suppress the hydrogen emission (Roth et al. 2016). Thus, the line
ratio is an observable helpful to discuss the formation of the compact disk, and further its time evolution• ToO Access :will tell
us whether the accretion flow evolves between a compact disk and an extended flow.

Our on-going project using Seimei/KOOLS-IFU

3

Only ToO access for ``spectroscopy” is invaluable
to study time-variable
TDEs
Observation
Details

To determine
the requested: nights (i.e., one night for each TDE,
• z < 0.2
• Source Identificationand threetotal),
we referred
to Khabibullin
et foral. (2014),• log(MBH/Msun) > 6
FoVnights
of ~ 8in
arcseccoversthe positional
uncertaintylog(Mgal/Msun) > 9.5
who studied
howsources
far awayby eROSITA.
eROSITA(i.e.,TDEs
d ~ 9 could
arcsecbe
fordetected~ 3 sigma)
Following one of their results, we considered SMBHs with 106 Msun
• High Sensitivity
from which
eROSITA :can find TDEs up to z ∼ 0.2. Given an
Suﬃcient
spectroscopy
sensitivity
for mostofobserved
ratio between
the SMBH
mass andthe galaxy
stellar mass
eROSITAcandidates
at z host< 0.2galaxies would be heavier(∼1/3000;ReinesTDE
et al.
2015), the→ for 18.5 mag. sources, 1 hr for 15 sigma detection per 10Å
than ∼ 3 × 109 Msun . Figure 3 shows the r-band mag distribution→ easier for brighter optical TDE flaresof the corresponding galaxies, created with a HSC galaxy catalog.
(※ from HSC data)
A peak can be seen at 18.5 mag. Then, the requested nights areadjusted so that we can detect H↵ emission from candidates down
r-bandmagnitude
Fig. 3: r-band
magnitude
distribution of
to the level of the host galaxy (= 18.5 mag). H↵ emission line isOptical studies of (optically selected) TDE candidates so far
U NIFIED TDE MODEL

3

ou
tfl
ow

f

:
w ng
flo ssiut ocet o repr
s
a ic
t
ba

ia
ad

slow o:utflow essingion
onizatphotoidebrisdebrisstream
streamdisk

X-ray &
opticaldisk

reprocOpticaldebris stream

ou
tfl
ow

ow
tfl
ou

jet

• What is the size of accreting material ?
→ related to a big question on how to extract
angular momentum→ Reﬀ ~ [L/4πσTeﬀ4]0.5

mis
sion

fun
nel:
in

• There are various questions to be addressed

trin
sic e

jet

X-ray

Figure 1. A schematic picture showing the viewing-angle dependence for the observed emission from a TDE super-Eddington disk. The

emission from the inner disk is reprocessed by the optically thick outflows and outer disk. Only when the observer is looking into the opticallyThe
Astrophysical Journal, 842:29 (15pp), 2017 June 10
thin funnel is the inner disk exposed, which can reveal strong, beamed X-ray and EUV radiation. Otherwise, X-ray is reprocessed into

Hung et al.

optical/NUV emission via photoionization (in denser outflow or disk at high inclination angles) or adiabatic cooling (in ultrafast outflow at low
inclination angles). A jet is included in the picture for completeness, though most TDEs may not produce jets.

used a spherically symmetric envelope with ad-hoc profile
for calculations. A general relativistic simulation of a superEddington TDE disk is the key to provide a clear, qualitative
understanding of the outflow profile and the viewing-angle
dependence of the observed emission.
2.1. Fully 3D GR Radiation MHD simulation set up
We simulate a super-Eddington TDE disk using thefully 3D general relativistic radiation magnetohydrodynamics (GRRMHD) code called HARMRAD (McKinney et al.
2014), which treats radiation under M1 closure (Levermore
1984). The gas is assumed to have solar chemical abundances (mass fractions of H, He, and “metals”, respectively,
X = 0.7,Y = 0.28, Z = 0.02), so its electron scattering opacityis es = 0.2(1 + X) cm2 g-1 . For absorption and emission,
frequency-mean opacities are used (see McKinney et al.
(2015) for the expression; also the Chianti opacity is turned

Figure 6. Time evolution of iPTF16axa. Upper: the time evolution ofoff as it is unimportant for the TDE disk temperature). Therblackbody
temperature of iPTF16axa. The black dotted line marks the mean
mal Comptonization is also included. There are
128 cells inradius r from 1.2 R to 10 R (R = GM /c is the gravof 3×104 K. Middle: the evolution of integrated UV–optical luminosity.
T increasing
itational radius of the black hole) with cell size bb
exponentially until r ⇡ 500 R and then even faster,
The64 cells
black dotted line shows the t−5/3 prediction from the light curve with ain ✓-grid from 0 to ⇡ with finer resolution in the jet and disk
region, and 32 cells in -grid spanning uniformly
from luminosity0
peak
indicated in Figure 4. Total power emitted (area under theto 2⇡. We provide large-scale poloidal magnetic flux to thedotted line integrated from tdisc to td ) is 5.5×1050 erg. Lower: the time
initial disk with plasma beta of ⇠ 20 - 30.
We tailor the parameters and initial conditions to be studyof the blackbody radius inferred from SED ﬁtting.
TDE disks. The SMBH has a mass M⇡ evolution
5 ⇥ 10 M
g

5

g

g

BH

2

g

BH

6

and spin parameter a = 0.8. The initial disk is Keplerian witha rest-mass density that is Gaussian in angle with a heightto-radius ratio of H/R ⇡ 0.3. Radially the density follows apower-law of ⇢ / r-1.3 out to Rdisk = 500 Rg and exponentially decays with r > Rdisk . We set up a small disk, since in
TDEs stellar debris is provided to the SMBH from a close

MJD
57551.2357556.82
57571.6357576.48
57585.93
57587.5957592.58
57594.42
57599.09
57608.057613.24
57620.54

Table 1
Blackbody Fitting from Light Curvest−tdiscdays

BB Temperature
104 K

BB Radius
1014 cm

Luminosity
1043erg s−1

14
19
34
39
49
50
55
57
62
71
76
83

3.13±0.23
2.82±0.22.64±0.08
3.14±0.24
2.82±0.19
3.15±0.11
2.8±0.2
2.95±0.12
2.66±0.29
3.91±0.68
3.08±0.34
2.88±0.35

3.98±2.27
4.1±2.31
3.59±1.86
2.6±1.49
3.24±1.81
2.68±1.42
2.8±1.57
2.94±1.57
3.19±1.93
1.76±1.32
2.27±1.41
2.35±1.47

10.91±1.14
7.56±0.66
4.44±0.13
4.67±0.51
4.71±0.37
5.02±0.28
3.46±0.29
4.67±0.25
3.63±0.45
5.11±1.69
3.32±0.53
2.68±0.43

We select the regions outside of the Balmer lines and the He II
emission line in the host-subtracted spectrum to estimate thecontinuum. The line-free regions are ﬁt with a ﬁfth-order Legendre
Figure 5. Blackbody ﬁt for the transient SEDs. The time indicated in each panel shows the time elapsed since discovery tdisc=MJD 57537.4. The blackbody
temperature remains roughly constant with a mean temperature 3×104 K over time.

We attempted fitting the observed H↵ emission line
profile on day-51 with the elliptical disk model by minimizing 2 in the following parameter space: 1 < q < 4,
100 < disk < 1, 000 km s 1 , 0  i < ⇡2 rad, 0 <e < 1, 0 < 2⇡ rad, 100 < ⇠1 < 3, 000 Rg , and100 < ⇠2 < 10, 000 Rg . We show the model parameters and the ranges of the flat priors in Table 3. We
noticed that the best-fitting disk-only model does not
capture the sharp peaks around the edges of the profile, but instead has broader peaks in order to matchthe high central flux. Alternatively, if we attempt tofit only the edges of the profile by excluding the central
data (with | v| < 3000 km s 1 ), the best-fitting model,
which is tabulated in Table 4, then lacks the flux seenat the center of the line profile (shown as the green line
in Figure 6). Clearly, a second component is needed toaccount for the central flux.
In order to match the full line profile, we add a Gaussian component to the disk model. We smoothed thedata of the H↵ emission with a boxcar function with alength of ⇠5 Å to erase fine structures that cannot be
captured by the emission-line model. Using the emcee
package (Foreman-Mackey et al. 2013), we sampled theposterior distribution of a total of ten parameters (theseven disk parameters listed above and an additional
three parameters that describe the Gaussian: amplitude, center (µ0 ; velocity oﬀset from the systemic velocity of the host galaxy), and width ( g ; standard deviation) using flat priors over the aforementioned ranges
for the disk parameters and over the following ranges for
the Gaussian component: 3, 000 < µ0 < 3, 000 km s 1 ,and 0 < g < 20, 000 km s 1 (Table 3). The model
parameter g can be translated to the FWHM
p of theGaussian profile via the relation FWHM= 2 2log2 g .
We normalized our flux to range between 0 and 1 and use
this same range as the flat prior for the amplitude of theGaussian. The posterior distributions of the model parameters are shown in Figure 7. The MCMC approachprovides a more general description of the model parameters, instead of forcing them to be Gaussian distributed
as in traditional fitting methods. In addition, we canalso visualize the correlation between diﬀerent parameters in the joint distributions, such as that between i,

Optical studies of (optically selected) TDE candidates so far(Hung+20)• There are various questions to be addressed

Disk

Gaussian(outflow?)• What is the size of accreting material ?
→ related to a big question on how to extract
angular momentum→ Reﬀ ~ [L/4πσTeﬀ4]0.5(Holoien+16)

f ,cont )Figure 6. t = 51-day spectrum zoomed in on the H↵ emission line profile and best-fitting model. The top and bottom
1.0
926
panels are of the same flux scale.
The black histogram in the430
top 0.8
panel shows the binned spectrum while the grey curve
marks the original spectrum with host light subtracted. The
red 0.6
curves show the best-fitting two-component model derived from the H↵ profile (dashed box; enlarged in lower
0.4
panel), rescaled by 1, 0.82, and 0.62 and shifted to the rest
wavelengthof H↵, H , and H emission, respectively, on top
0.2
of the local continuum. The local continuum is estimated by
0.0
linearlyinterpolating the regions bracketing each emission
line. The black curve in the lower panel shows the zoom-in
1.0of the continuum-subtractedview
H↵ line profile. The solid
85
145
73
132
red curve is the best-fitting two-component(disk + Gaus0.8sian) model with the parameters indicated in Table 4. The
two0.6components are shown separately in orange dotted line
(Gaussian) and in blue dashed line (disk). The green line
0.4 the best-fitting model with only the disk componentshows
by omitting the center part of the line flux. It is clear that0.2
the disk-only model cannot reproduce the all flux near zero
velocity.
0.0

Normalized (f• What is the structure of accreting matter ?
→ Disk and outflow→ Evolution of broad lines (narrower with brighter)→ Reverberation mapping b/w X-ray and optical

The Nearby T

-2

-1

0

1

2

-1

0

1

2

we note that the relative
strength of the two peaks isv(H↵)/104 (km s 1)
orientation-dependent in elliptical disks.
To 8.
summarize,
thetheday-51
spectrumis bestFigureEvolution ofH↵ (leftpanel) andHe IIdescribed
4686 Å (right panel) line profileWe attempted fitting the observed H↵ emission line
profile on day-51 with the elliptical disk model by minimizing 2 in the following parameter space: 1 < q < 4,
100 < disk < 1, 000 km s 1 , 0  i < ⇡2 rad, 0 <e < 1, 0 < 2⇡ rad, 100 < ⇠1 < 3, 000 Rg , and100 < ⇠2 < 10, 000 Rg . We show the model parameters and the ranges of the flat priors in Table 3. We
noticed that the best-fitting disk-only model does not
capture the sharp peaks around the edges of the profile, but instead has broader peaks in order to matchthe high central flux. Alternatively, if we attempt tofit only the edges of the profile by excluding the central
data (with | v| < 3000 km s 1 ), the best-fitting model,
which is tabulated in Table 4, then lacks the flux seenat the center of the line profile (shown as the green line
in Figure 6). Clearly, a second component is needed toaccount for the central flux.
In order to match the full line profile, we add a Gaussian component to the disk model. We smoothed thedata of the H↵ emission with a boxcar function with alength of ⇠5 Å to erase fine structures that cannot be
captured by the emission-line model. Using the emcee
package (Foreman-Mackey et al. 2013), we sampled theposterior distribution of a total of ten parameters (theseven disk parameters listed above and an additional
three parameters that describe the Gaussian: amplitude, center (µ0 ; velocity oﬀset from the systemic velocity of the host galaxy), and width ( g ; standard deviation) using flat priors over the aforementioned ranges
for the disk parameters and over the following ranges for
the Gaussian component: 3, 000 < µ0 < 3, 000 km s 1 ,and 0 < g < 20, 000 km s 1 (Table 3). The model
parameter g can be translated to the FWHM
p of theGaussian profile via the relation FWHM= 2 2log2 g .
We normalized our flux to range between 0 and 1 and use
this same range as the flat prior for the amplitude of theGaussian. The posterior distributions of the model parameters are shown in Figure 7. The MCMC approachprovides a more general description of the model parameters, instead of forcing them to be Gaussian distributed
as in traditional fitting methods. In addition, we canalso visualize the correlation between diﬀerent parameters in the joint distributions, such as that between i,

Optical studies of (optically selected) TDE candidates so far(Hung+20)• There are various questions to be addressed

Disk

Gaussian(outflow?)• What is the size of accreting material ?
→ related to a big question on how to extract
angular momentum→ Reﬀ ~ [L/4πσTeﬀ4]0.5• How are all of these aspects aﬀected by the propertiesof the BH and those of the disrupted star ?
→ need a large sample(Holoien+16)

f ,cont )Figure 6. t = 51-day spectrum zoomed in on the H↵ emission line profile and best-fitting model. The top and bottom
1.0
926
panels are of the same flux scale.
The black histogram in the430
top 0.8
panel shows the binned spectrum while the grey curve
marks the original spectrum with host light subtracted. The
red 0.6
curves show the best-fitting two-component model derived from the H↵ profile (dashed box; enlarged in lower
0.4
panel), rescaled by 1, 0.82, and 0.62 and shifted to the rest
wavelengthof H↵, H , and H emission, respectively, on top
0.2
of the local continuum. The local continuum is estimated by
0.0
linearlyinterpolating the regions bracketing each emission
line. The black curve in the lower panel shows the zoom-in
1.0of the continuum-subtractedview
H↵ line profile. The solid
85
145
73
132
red curve is the best-fitting two-component(disk + Gaus0.8sian) model with the parameters indicated in Table 4. The
two0.6components are shown separately in orange dotted line
(Gaussian) and in blue dashed line (disk). The green line
0.4 the best-fitting model with only the disk componentshows
by omitting the center part of the line flux. It is clear that0.2
the disk-only model cannot reproduce the all flux near zero
velocity.
0.0

Normalized (f• What is the structure of accreting matter ?
→ Disk and outflow→ Evolution of broad lines (narrower with brighter)→ Reverberation mapping b/w X-ray and optical

The Nearby T

-2

-1

0

1

2

-1

0

1

2

we note that the relative
strength of the two peaks isv(H↵)/104 (km s 1)
orientation-dependent in elliptical disks.
To 8.
summarize,
thetheday-51
spectrumis bestFigureEvolution ofH↵ (leftpanel) andHe IIdescribed
4686 Å (right panel) line profile

Actual optical obs. of
three X-ray Events(Only interesting objects are selected)

Source 1: eRASSt J093641.9+111347- Ultra-soft event in a massive galaxy (Grotova et al. 2021, ATel#14668)
eRASS2eRASS3

kT~70eV• Very soft (kT~70eV) discovery on May 13th 2021
• f0.2-2.3 keV~ 8.7x10-13 cgs, i.e. ~13x times brighter than eRASS:2
3-sigma UL• L0.2-2.3keV~3x1043 cgs• NICER (PI: Z. Liu) & Swift follow-up (PI: I. Grotova) indicates source isfading now. Potential hardening observed• No UV brightening detected• WISE color W1-W2=0.15mag

Source 1: J0936 - no optical change
Seimei: 27th May (= 2 weeks later obs.)

SDSS• Host is FR II radio galaxy at z=0.12• Archival SDSS spectrum shows no emission
lines• MSMBH~108.8 MSun (Capetti et al. 2017)• Optical spectrum ~15d post eROSITA
discovery suggests no significant change(Observations: S. Ciroi & F. di Mille)

TDEs (cand.) are not always bright at opt. !
→ Comprehensive understanding of TDEs
needs a study of X-ray TDEs

Source 2: eROSITA_DE’s brightest nuclear transient (up to now)(Fig.: D. Homan)(Fig.: A. Malyali)• Brightest event in eROSITA_DE so far(f0.2-5keV~1.4x10-11 cgs)
• Very soft X-ray spectrum (Gamma~5, kT~80eV)• L0.2-5keV,peak~5x1044 cgs (z=0.1)• (Also detected by GAIA & ATLAS)• Archival WISE color (W1-W2=0.35mag) does notsuggest strong previous AGN activity

Source 2: TDE or AGN ingnition? (Honest answer: we don’t know, yet.)• Host galaxy member of a group at z=0.1• broad Balmer lines with narrow components• No significant spectral change observed (until June 2021)
GROND(Fig.: D. Homan; Observations: M. Krumpe, D. Buckley, S. Ciroi, F. di Mille)

2020 December

2021 June 2nd(~ 6 months later)~ 1e44 erg/s(Fig. & Observations: S. Ciroi & F. diMille)~ 5e43 erg/sSource 3: AT 2019avd, extreme X-ray flaring and optical variability from a quiescent galaxy
A. Malyali et al. (MPE, A&A 2021)(Fig.: A. Malyali)(Fig.: A. Malyali)

IR

eROSITA
ZTF• eRASS1 discovery vs archival XMM observations• X-ray properties (eROSITA, Swift XRT) suggest a
TDE candidate• However, optical/mid-IR light curve inconsistent
with standard TDE interpretation• A new class of nuclear transients?
Stellar binary TDE? TDE involving a SMBHB?

Optical

X-ray
Brightens by factor of >600 in 0.2-2 keV
relative to 2015 XMM upper limit

Source 3: AT 2019avd, extreme X-ray flaring and optical variability from a quiescent galaxy
A. Malyali et al. (MPE, A&A 2021)• Detection of Bowen fluorescence lines of NIII 4640Å
(i.e., Strong UV/X-ray photos → HeII → OIII → NIII)• The presence of UV/X-ray absorbers• A high ratio of F(NIII 4640Å)/F(Hβ) suggests
the presence of high density gas (nH > 109.5 cm-3), comparable to those of BLRs of AGNs

Evolution b/w 14 months

Summary•~30 X-ray soft nuclear transients, or TDE candidates, per eRASS (~ 6 months),
which will be until ~ 2024 → Now is the time for X-ray TDE studies•Optical data is very important for understanding the mass accretion in TDEs
•Optical spectroscopy have unveiled optical properties diﬀerent from thosefound for optically selected TDE candidates•Non-canonical behavior
i.e., ‘If you have seen one TDE candidate, you have seen one TDE candidate’