# Transient fading X-ray emission detected during the optical rise of a tidal
disruption event

A. Malyali1, A. Rau1, C. Bonnerot2, A. J. Goodwin3, Z. Liu1, G. E. Anderson3,
J. Brink4,5, D. A. H. Buckley4,5,6, A. Merloni1, J. C. A. Miller-Jones3, I.
Grotova1, A. Kawka3
1Max-Planck-Institut für extraterrestrische Physik, Giessenbachstrasse 1,
85748 Garching, Germany
2School of Physics and Astronomy & Institute for Gravitational Wave Astronomy,
University of Birmingham, Birmingham B15 2TT, UK
3International Centre for Radio Astronomy Research, Curtin University, GPO Box
U1987, Perth, WA 6845, Australia
4South African Astronomical Observatory, PO Box 9, Observatory Rd, 7935
Observatory, Cape Town, South Africa
5Department of Astronomy, University of Cape Town, Private Bag X3, Rondebosch
7701, South Africa
6Department of Physics, University of the Free State, PO Box 339, Bloemfontein
9300, South Africa E-mail<EMAIL_ADDRESS>
(Accepted XXX. Received YYY; in original form ZZZ)

###### Abstract

We report on the SRG/eROSITA detection of ultra-soft ($kT=47^{+5}_{-5}$ eV)
X-ray emission ($L_{\mathrm{X}}=$$2.5^{+0.6}_{-0.5}\times 10^{43}$ erg s -1)
from the tidal disruption event (TDE) candidate AT 2022dsb $\sim$14 days
before peak optical brightness. As the optical luminosity increases after the
eROSITA detection, then the 0.2–2 keV observed flux decays, decreasing by a
factor of $\sim$39 over the 19 days after the initial X-ray detection. Multi-
epoch optical spectroscopic follow-up observations reveal transient broad
Balmer emission lines and a broad He ii 4686Å emission complex with respect to
the pre-outburst spectrum. Despite the early drop in the observed X-ray flux,
the He ii 4686Å complex is still detected for $\sim$40 days after the optical
peak, suggesting the persistence of an obscured, hard ionising source in the
system. Three outflow signatures are also detected at early times: i)
blueshifted H$\alpha$ emission lines in a pre-peak optical spectrum, ii)
transient radio emission, and iii) blueshifted Ly$\alpha$ absorption lines.
The joint evolution of this early-time X-ray emission, the He ii 4686Å complex
and these outflow signatures suggests that the X-ray emitting disc (formed
promptly in this TDE) is still present after optical peak, but may have been
enshrouded by optically thick debris, leading to the X-ray faintness in the
months after the disruption. If the observed early-time properties in this TDE
are not unique to this system, then other TDEs may also be X-ray bright at
early times and become X-ray faint upon being veiled by debris launched
shortly after the onset of circularisation.

###### keywords:

accretion, accretion discs – galaxies: nuclei – black hole physics –
transients: tidal disruption events –

††pubyear: 2023††pagerange: Transient fading X-ray emission detected during
the optical rise of a tidal disruption event–B

##  1 Introduction

The number of stellar tidal disruption event (TDE) candidates identified in
recent years has greatly increased, largely fuelled by the increasing number
of wide-field, high-cadence time-domain surveys operating across the
electromagnetic spectrum. Although early theoretical work predicted TDEs to
produce large amplitude, ultra-soft X-ray flares originating from the centres
of galaxies (Rees, 1988) – consistent with the first TDE candidates identified
by ROSAT (Trümper, 1982) in the 1990s (Bade et al., 1996; Grupe et al., 1999;
Komossa & Greiner, 1999; Komossa & Bade, 1999; Greiner et al., 2000) – the
majority of optically-selected TDE candidates do not show transient X-ray
emission (van Velzen et al., 2021; Hammerstein et al., 2023). To explain the
dearth of X-rays in these systems, it has been suggested that the optical
emission is produced by the debris circularisation process instead of
accretion (stream-stream collisions; Piran et al. 2015; Shiokawa et al. 2015),
or that a large fraction of the X-ray emission is reprocessed to optical/UV
bands by debris enveloping the nascent disc (Loeb & Ulmer, 1997; Lodato &
Rossi, 2011; Miller, 2015; Metzger & Stone, 2016; Dai et al., 2018; Lu &
Bonnerot, 2020).

Optical spectroscopic follow-up of optically bright TDEs has led to the
classifications of TDEs into different spectral types (e.g. Arcavi et al.
2014; Leloudas et al. 2019; van Velzen et al. 2021), depending on the emission
lines seen in the spectra. These are i) ‘H’, which show transient broad Balmer
emission lines, ii) ‘H+He’, showing transient broad Balmer emission lines and
a broad emission complex around He ii 4686A, and iii) ‘He’, which show a
transient broad He ii 4686A emission feature but no Balmer emission. An
additional spectral TDE class not common in recent optically-selected TDE
samples are the extreme coronal line emitters (ECLEs; Komossa et al. 2009;
Wang et al. 2011, 2012), which show strong emission from high-ionisation
coronal lines with respect to their narrow [O iii] 5007 Å emission. A hard
ionising source (photons with energy above 54 eV) is needed to produce the He
ii emission seen in ‘He’ and ‘H+He’ TDEs (herein collectively referred to as
He-TDEs), yet the majority of TDE candidates even in these classes do not show
transient X-ray emission (Hammerstein et al., 2023). As it is thought that
these hard photons originate from the high-energy tail of the newly formed
disc, then the combination of the X-ray faintness and the He ii emission in
these systems has been suggested as evidence for ‘obscured accretion’
(Leloudas et al., 2019), where an accretion disc has formed in these systems,
but its high-energy emission gets reprocessed into the optical band by an
optically-thick gaseous envelope. Several TDE candidates have also shown broad
He ii lines close to peak optical brightness (Blagorodnova et al., 2017;
Nicholl et al., 2020; Wevers et al., 2022), which under the assumption of an
obscured accretion-driven origin, suggests efficient circularisation of the
debris into a disc post-disruption.

Here, we report on multi-wavelength observations of the TDE candidate AT
2022dsb, which shows a factor of 39 decrease in its 0.2–2 keV observed flux
during the optical rise. Section 2 describes the discovery of AT 2022dsb,
whilst sections 3 and 4 detail multi-wavelength observations of the system and
their analysis, respectively. In section 5, we review previous X-ray
observations of TDEs at early times and compare these with AT 2022dsb. The
implications of our observational campaign are discussed in section 6, and our
conclusions in section 7. All magnitudes are reported in the AB system and
corrected for Galactic extinction using $A_{\mathrm{V}}=0.62$ mag, obtained
from Schlafly & Finkbeiner (2011), $R_{\mathrm{V}}=3.1$ and a Cardelli
extinction law (Cardelli et al., 1989). The effective wavelength for each
filter was retrieved from the SVO Filter Profile
Service111http://svo2.cab.inta-csic.es/theory/fps/. All dates and times will
be reported in universal time (UT).

##  2 Discovery

AT 2022dsb/ eRASSt J154221.6-224012 was independently discovered by the
extended ROentgen Survey with an Imaging Telescope Array (eROSITA; Predehl et
al. 2021), the soft X-ray instrument on board the Spektrum-Roentgen-Gamma
(SRG; Sunyaev et al. 2021) observatory, during a systematic search for TDE
candidates in the fifth eROSITA All-Sky Survey (eRASS5), when it was observed
on 2022-02-17 as a new, bright (0.2-2 keV observed flux of $\sim 3\times
10^{-13}$ erg cm-2 s-1), ultra-soft (section 4.1) X-ray point source. Using
the eROSITA Science Analysis Software pipeline (eSASS222Version:
eSASSusers_211214.; Brunner et al. 2022), the source was localised to
(RAJ2000, Dec${}_{\mathrm{J2000}})$=(15h42m21.6s,-22∘40′12.1′′), with a
1$\sigma$ positional uncertainty of 1.9′′ (68% confidence), consistent with
the galaxy ESO 583-G004 at $z=0.0235$ (Fig. 1). No X-ray source had been
detected within 30″of this position in any of the previous four eRASS, with a
3$\sigma$ upper limit on the 0.2-2 keV band flux of $5\times 10^{-14}$ erg
cm-2 s-1, assuming the same spectral model fitted to the eRASS5 spectrum (Liu
et al., 2022). The last non-detection by eROSITA occurred $\sim$6 months
before the eRASS5 detection.

AT 2022dsb was later publicly classified on 2022-03-02 as a TDE candidate in
the TNS report TNSCR-2022-584 (Fulton et al., 2022), after the discovery and
reporting of optical transient emission (associated to the nucleus of the host
galaxy ESO 583-G004) initially by ASAS-SN on 2022-03-01 in Stanek & Kochanek
(2022), and then by both the Asteroid Terrestrial Impact Last Alert System
(ATLAS; Tonry et al. 2018), and the Zwicky Transient Facility (ZTF; Bellm et
al. 2019; Graham et al. 2019) by the ALeRCE alert broker (Förster et al.,
2021).

Figure 1: Finder chart for AT 2022dsb (DESI LS DR10 $g$-band image). The red
circle denotes the 3$\sigma$ uncertainty on the eROSITA source position in
eRASS5, whilst the dark orange star marks the Gaia DR3 optical centre of the
host galaxy ESO 583-G004.

###  2.1 Host galaxy

A pre-outburst optical spectrum of ESO 583-G004 was taken in 2002 during the
Six-Degree Field (6dF; Jones et al. 2009) galaxy survey. A recent analysis of
the narrow emission lines in this optical spectrum classified the system as a
type II AGN (Chen et al., 2022), according to the criteria presented in Kewley
et al. (2001). However, the pre-outburst AllWISE (Wright et al., 2010; Mainzer
et al., 2014) colour of the host, $W1-W2=0.00\pm 0.03$ mag, suggests that its
mid-infrared emission is dominated by the galaxy light, instead of the
luminous emission from a dusty torus surrounding an AGN (Stern et al., 2012;
Assef et al., 2018). ESO 583-G004 may have hosted a low-luminosity AGN prior
to the outburst of AT 2022dsb, similar to other TDE candidates (e.g.
ASASSN-14li, Holoien et al. 2016; AT 2019qiz, Nicholl et al. 2020).

We fitted the DESI Legacy DR10 (Dey et al., 2019) archival photometry of the
host galaxy ($g$, $i$, $W1$, $W2$, $W3$ and $W4$ bands333No $z$-band
photometry was available for the host of AT 2022dsb in DR10, and all DR10
photometry was taken before August 2021.) with the stellar population
inference tool Prospector (Johnson et al., 2021), which uses a python wrapper
(Foreman-Mackey et al., 2014) of the Flexible Stellar Population Synthesis
code (Conroy & Gunn, 2010) for generating the SEDs of stellar populations. The
SED model includes both stellar and nebular emission, as well as dust
attenuation and emission, and adopts a Chabrier initial mass function (IMF;
Chabrier 2003); the free parameters are the total stellar mass of the galaxy
($M_{\star}$, the sum of both living and remnant stars), the metallicity
($\log(Z/Z_{\odot})$), the age of the galaxy ($t_{\mathrm{age}}$), the decay
timescale under an exponentially declining star formation model
($\tau_{\mathrm{SF}}$), and the host galaxy dust extinction
($A_{\mathrm{V}}$). Posterior distributions were sampled using the dynamic
nested sampler (Skilling, 2004, 2006; Higson et al., 2019) dynesty (Speagle,
2020; Koposov et al., 2023), with the posterior model shown in Fig. 2 and the
parameter estimates in Table 1. From the inferred $M_{\star}=7^{+4}_{-3}\times
10^{10}\,\rm{M_{\odot}}$and using the relation between $M_{\mathrm{BH}}$ and
$M_{\star}$ in Reines & Volonteri (2015), we infer
$\log[M_{\mathrm{BH}}/M_{\odot}]=7.3^{+0.2}_{-0.3}$.

Figure 2: SED fit to the LS DR10 photometry of the host galaxy of AT 2022dsb. The observed photometry is plotted with black edged circular markers, whilst the posterior model and model photometry is shown in blue (solid line represents the median, shaded region encloses the middle 90% of the posterior) and red, respectively.  Table 1: Host galaxy properties inferred via SED fitting to the archival photometry, with CR denoting the credible region for a parameter. Parameter | 68% CR
---|---
$M_{\mathrm{\star}}$$/10^{10}M_{\odot}$ | $6.9^{+3.9}_{-3.3}$
$\log[Z/\mathrm{Z}_{\odot}]$ | $-0.7^{+0.3}_{-0.3}$
$A_{\mathrm{V}}$/mag | $0.62^{+0.05}_{-0.06}$
$t_{\mathrm{age}}$/Gyr | $5.4^{+4.3}_{-3.2}$
$\tau_{\mathrm{SF}}$/Gyr | $0.2^{+0.5}_{-0.1}$

##  3 Observations and data reduction

After the initial eROSITA discovery, additional X-ray observations of AT
2022dsb were obtained with XMM-Newton (section 3.1.2) and Swift XRT (section
3.1.3); collectively, these sample the X-ray emission from a TDE pre-peak,
near-peak and post-peak optical brightness. The optical, UV and radio
evolution was also monitored with ground-based photometry (section 3.2.1),
Swift UVOT (section 3.2.2) and the Australia Telescope Compact Array (ATCA;
section 3.4), respectively. The full multi-wavelength lightcurve of AT 2022dsb
is depicted in Fig. 3, and a comparison of the optical lightcurve with other
TDE candidates is shown in Fig. 4. A log of all X-ray observations and the
inferred fluxes is presented in Table 2, whilst the optical and UV photometry
can be found in Table 9.

Figure 3: 0.2–2 keV X-ray (top) and optical-UV (bottom) evolution of AT 2022dsb. Solid markers denote $>3\sigma$ detections in each epoch, whereas translucent triangles mark $3\sigma$ upper limits. The vertical red line marks the eROSITA observation of AT 2022dsb, whilst the vertical orange line marks the inferred time of optical peak on MJD 59641 (section 4.2), $\sim 14$ days after the eRASS5 detection.  Figure 4: Comparison of the ATLAS $o$-band lightcurve of AT 2022dsb (red stars) with the $g$-band lightcurves of ZTF-selected TDEs (blue markers) reported in Hammerstein et al. (2023). TDEs of similar peak absolute magnitudes are highlighted in non-dark blue colours, and we include data from the ‘faint and fast’ TDE iPTF-16fnl (Blagorodnova et al., 2017), and the ‘faint and slow’ TDE eRASSt J074426.3+291606 (the faintest optically-bright TDE observed to date, J0744; Malyali et al. 2023b). Table 2: X-ray lightcurve of AT 2022dsb. $F_{\rm 0.2-2keV,obs}$ and $F_{\rm 0.2-2keV,unabs}$ are the observed (not corrected for Galactic absorption) and unabsorbed 0.2–2 keV band fluxes in units of erg cm-2 s-1. $\log[L_{\rm 0.2-2keV}$] is inferred from $F_{\rm 0.2-2keV,unabs}$. MJD is computed from the midpoint of MJDstart and MJDstop. The fluxes have been estimated from the best fitting model (Table 4), with the 3$\sigma$ upper limits on the count rates converted to fluxes using the best fitting model to the eRASS5 spectrum. MJD | MJDstart | MJDstop | Instrument | ObsID | $\log[F_{\rm 0.2-2keV,obs}]$ |  $\log[F_{\rm 0.2-2keV,unabs}$] |  $\log[L_{\rm 0.2-2keV}$]
---|---|---|---|---|---|---|---
59627.939 | 59627.439 | 59628.439 | eROSITA | eRASS5 | $-12.46^{+0.07}_{-0.08}$ | $-10.79^{+0.10}_{-0.14}$ | $43.34^{+0.10}_{-0.14}$
59643.115 | 59643.004 | 59643.226 | XRT | 00015054002 |  $<$-12.69 |  $<$-11.11 |  $<$43.02
59646.226 | 59646.088 | 59646.363 | EPIC PN | XMM1 | $-14.05^{+0.05}_{-0.06}$ | $-13.50^{+0.10}_{-0.10}$ | $40.62^{+0.10}_{-0.10}$
59649.772 | 59649.702 | 59649.842 | XRT | 00015054003 |  $<$-12.71 |  $<$-11.12 |  $<$43.01
59656.156 | 59656.150 | 59656.161 | XRT | 00015054005 |  $<$-12.62 |  $<$-11.03 |  $<$43.10
59663.969 | 59663.968 | 59663.969 | XRT | 00015054006 |  $<$-10.81 |  $<$-9.22 |  $<$44.90
59668.128 | 59668.030 | 59668.225 | XRT | 00015054007 |  $<$-12.65 |  $<$-11.07 |  $<$43.06
59673.808 | 59673.805 | 59673.811 | XRT | 00015054008 |  $<$-12.26 |  $<$-10.67 |  $<$43.45
59677.316 | 59677.178 | 59677.453 | XRT | 00015054009 |  $<$-12.12 |  $<$-10.53 |  $<$43.60
59683.209 | 59683.205 | 59683.214 | XRT | 00015054010 |  $<$-12.38 |  $<$-10.80 |  $<$43.33
59696.480 | 59696.206 | 59696.755 | XRT | 00015054011 |  $<$-12.53 |  $<$-10.94 |  $<$43.19
59701.425 | 59701.251 | 59701.598 | XRT | 00015054012 |  $<$-12.55 |  $<$-10.96 |  $<$43.16
59704.749 | 59704.513 | 59704.986 | XRT | 00015054013 |  $<$-12.51 |  $<$-10.93 |  $<$43.20
59729.976 | 59729.051 | 59730.902 | XRT | 00015054014 |  $<$-12.27 |  $<$-10.68 |  $<$43.44
59735.808 | 59735.801 | 59735.816 | XRT | 00015054015 |  $<$-12.39 |  $<$-10.80 |  $<$43.32
59820.520 | 59820.376 | 59820.664 | EPIC PN | XMM2 | $-14.18^{+0.08}_{-0.08}$ | $-13.38^{+0.16}_{-0.15}$ | $40.75^{+0.16}_{-0.15}$
59855.517 | 59855.146 | 59855.888 | XRT | 00015054016 |  $<$-12.91 |  $<$-11.32 |  $<$42.81

###  3.1 X-ray

####  3.1.1 eROSITA

The position of AT 2022dsb was observed by eROSITA during the first four
eRASS, (denoted eRASS1, eRASS2, eRASS3, and eRASS4, respectively) on
2020-02-27, 2020-08-26, 2021-02-11, and 2021-08-20. During eRASS5, eROSITA
first observed AT 2022dsb on 2022-02-17, scanning over its position seven
times over the following day, with each visit separated by four hours (Fig.
5). Using the eSASS task SRCTOOL, we generated source and background spectra
by extracting counts from a source aperture of radius 30′′, and a background
annulus of inner and outer radii 90′′ and 240′′, respectively, with both
apertures centred on the eROSITA position of AT 2022dsb. The same source and
background apertures were used to generate a 0.2–2 keV lightcurve using
SRCTOOL, shown in Fig. 5. AT 2022dsb is clearly detected above background by
eROSITA in each of the seven observations within eRASS5 (i.e. is persistently
bright, instead of showing a ‘one-off’ short flaring), providing a lower limit
of 1 day on the duration of X-ray emission at early-times in AT 2022dsb. A log
of X-ray observations is presented in Table 2.

Figure 5: eRASS5 lightcurve of AT 2022dsb in the 0.2–2 keV band. The blue
markers denote the source count rates (corrected for vignetting), whilst the
grey markers show the estimated background count rates.
$t-t_{\mathrm{eRASS5,0}}$ is the time relative to the start of eROSITA’s
observations of AT 2022dsb in eRASS5 (MJD=59627.439). AT 2022dsb is
persistently bright over the day-long monitoring window in eRASS5.

####  3.1.2 XMM-Newton

Additional observations of AT 2022dsb were performed with XMM (P.I. Z. Liu),
with the first taking place $\sim$19 days after the eRASS5 detection on
2022-03-08, and then $\sim$173 days after this on 2022-08-29; observations
were carried out in imaging mode with the medium filter. To reduce and analyse
the XMM data, we used HEASOFT (version 6.29), the XMM Science Analysis
Software (SAS) (version 20211130_0941), and the latest calibration data files.
Calibrated event files were generated from the Observation Data Files (ODF)
using emproc and epproc for the MOS and PN cameras, respectively, and periods
of high particle background during each observation were filtered out
following the XMM Science Operation Centre recommended procedures. This
resulted in 18.0 ks and 13.3 ks exposures for the first and second
observations, respectively. Source spectra were extracted from a circle of
radius 20′′, centred on the Gaia EDR3 (Gaia Collaboration et al., 2021)
position of ESO 583-G004, whilst background spectra were extracted from an
annulus with inner and outer radii 76′′ and 144′′ respectively. Only events
with PATTERN<=4 and FLAG==0 were extracted for PN, whilst PATTERN<=12 was
applied for MOS1 and MOS2.

####  3.1.3 Swift XRT

AT 2022dsb was further monitored in the 0.3-10 keV band with the XRT
instrument (Burrows et al., 2005) on-board the Neil Gehrels Swift observatory
(Gehrels et al., 2004)444P.I. for Swift observations: A. Malyali, P.
Charalampopoulos, J. Hinkle, I. Lypova.. XRT observations commenced on
2022-03-05, $\sim$16 days after the eRASS5 observation, and were performed in
photon counting mode. These were then analysed with the online XRT product
building tool provided by the UK Swift Science Data Centre’s (UKSSDC) (Evans
et al., 2007, 2009). AT 2022dsb was not detected in any of the XRT
observations, with 3$\sigma$ upper limits on the 0.3-2 keV count rates
computed using the method presented in Kraft et al. (1991). These were then
converted to 0.2-2 keV fluxes using
webPIMMs555https://heasarc.gsfc.nasa.gov/cgi-bin/Tools/w3pimms/w3pimms.pl,
where we adopted the spectral model inferred from our BXA fit to the eRASS5
spectrum.

###  3.2 Photometry

####  3.2.1 Ground-based photometry

We obtained ATLAS $o$ and $c$ band (Tonry et al., 2018) lightcurves of AT
2022dsb using the online forced photometry server (Smith et al., 2020;
Shingles et al., 2021). For late-time observations (MJD$>$59635), we performed
a weighted rebin of the lightcurve into 1 day intervals. To improve the
sampling of the lightcurve around peak optical brightness, then no such
rebinning was performed for observations performed during the optical rise and
the early part of the decay ($59620<\mathrm{MJD}<59635$). To remove epochs of
low quality photometry in the ATLAS lightcurve, we discarded datapoints where
the semi-major axis of the fitted PSF model was greater than 3 pixels (1.86′′
per pixel).

In addition, $g$ and $r$-band lightcurves666This is generated from science
images that have already been reference image subtracted were generated using
the ZTF forced photometry service (Masci et al., 2019), which were then
calibrated using the method developed by Miller et al. (in prep.) for the ZTF
Bright Transient
Survey777https://github.com/BrightTransientSurvey/ztf_forced_phot. No
significant optical variability is seen in the ZTF lightcurves before the 2022
outburst, and we note that the ZTF observations do not sample the rise and
peak optical brightness (Fig. 3).

####  3.2.2 Swift UVOT

Over the course of the Swift monitoring campaign, AT 2022dsb was observed by
the UVOT (Roming et al., 2005) instrument across all filters ($V$, $B$, $U$,
$UVW1$, $UVM2$ and $UVW2$), although the number of filters used varied between
each observation (see photometry in Table 9). In this work, we use
observations performed only in the $UVW1$, $UVM2$ and $UVW2$ filters, since
the lightcurve sampling is highest in these bands, and the optical coverage is
already provided by ATLAS and ZTF. We first downloaded the level 2 UVOT sky
images from the UK Swift Science Data Centre, before computing aperture
photometry on these with the uvotsource task (HEASOFT v6.29, CALDB v20201215),
using a 5′′ radius source aperture, and a nearby, source-free circular
aperture of of radius 15′′ for the background. Lastly, the recommended Small
Scale Sensitivity
check888https://swift.gsfc.nasa.gov/analysis/uvot_digest/sss_check.html was
completed.

###  3.3 Optical spectroscopy

The first follow-up optical spectrum of AT 2022dsb was obtained on 2022-02-26
(MJD$=59636$, $\sim$5 days before optical peak), using the FLOYDS spectrograph
mounted on the 2m Las Cumbres Observatory (LCO; Brown et al. 2013)999Proposal
ID CON2022A-001, PI: M. Salvato. telescope at Haleakala Observatory. Data
processing and spectrum extraction were performed by the automatic FLOYDS
pipeline at LCO (further details on the spectroscopic data reduction are
presented in section B). This spectrum (Fig. 6) shows transient broad Balmer
emission lines (H$\alpha$, H$\beta$), a broad emission complex around 4600Å,
and a blue continuum, with the transient nature confirmed through comparison
to archival and late-time optical spectra. In addition, narrow emission lines
(H$\alpha$, H$\beta$, [N ii] 6548Å and 6583 Å, and the high-ionisation lines
[O iii] 4959Å and 5007 Å), as well as several host galaxy absorption features,
are clearly present. No strong blue continuum or broad emission lines were
seen in a pre-outburst optical spectrum taken on 2002-04-15 during the 6dF
Galaxy Survey (6dFGS; Jones et al. 2009). Over the 140 days of spectroscopic
monitoring of AT 2022dsb after optical peak, the strength of the broad
emission lines and the blue continuum relative to the host galaxy decreases
(Fig. 6). Zoom-in plots on the evolution of the H$\alpha$ and He ii complexes
are presented in Fig. 15.

Figure 6: Optical spectroscopic evolution of AT 2022dsb. The phase of the
observation with respect to the inferred optical peak
(MJD=$59640.9^{+0.5}_{-0.4}$) is shown on the right hand side above each
spectrum. Black, orange and red spectra were obtained using LCO/FLOYDS,
NTT/EFOSC2 and SALT/RSS, respectively. The archival spectrum of the host
galaxy is presented in Fig. 16.

###  3.4 Radio

####  3.4.1 Archival

The Karl G. Jansky Very Large Array Sky Survey (VLASS; Lacy et al. 2020)
observed the coordinates of AT2022dsb on 2020-11-03 and 2018-02-15,
approximately 1 and 4 years prior to the detection of the transient event.
There is no source present at the location of AT2022dsb in either of these
observations, with a 3$\sigma$ upper limit of 507$\mu$Jy and 419$\mu$Jy at 3
GHz for the 2020 and 2018 observations respectively.

####  3.4.2 Follow-up

We observed the coordinates of AT2022dsb on three occasions with the Australia
Telescope Compact Array (ATCA) between 2022 March and November (project C3334,
PI Anderson/Goodwin). Observations were taken in the 4-cm band with the dual
5.5 and 9 GHz receiver. Further, more detailed radio spectral monitoring of
AT2022dsb is being carried out and will be published in a follow-up paper
(Goodwin et al., in prep.). Because of the early eROSITA detection of AT
2022dsb, then the ATCA observations presented here represent one of the
earliest radio detections of a TDE.

The ATCA data were reduced using the Common Astronomy Software Application
(CASA v 5.6.3; The CASA Team et al. 2022) using standard procedures including
flux and bandpass calibration with PKS 1934-638 and phase calibration with PKS
1514-241. The target field was imaged using the CASA task tclean with an image
size of 4000 pixels and a cellsize of 0.3 arcsec at 5.5 GHz and an image size
of 4000 pixels and a cellsize of 0.2 arcsec at 9 GHz. In all observations a
point source was detected at the location of AT2022dsb. The flux density of
the point source was extracted in the image plane using the CASA task imfit by
fitting an elliptical Gaussian fixed to the size of the synthesized beam. A
summary of the ATCA observations is given in Table 3 and the 5.5 GHz
lightcurve of AT2022dsb is plotted in Figure 7 along with a selection of other
radio-detected TDEs for comparison. Both the variability of the detected 5.5
GHz and 9 GHz radio emission and that the initial detection is above the
3$\sigma$ VLASS 3 GHz upper limits years prior to the TDE suggest that the
radio emission is likely related to the transient event and is not purely host
galaxy emission. Although this VLASS upper limit is at a lower frequency than
the ATCA observations (5.5 and 9.0 GHz), the ATCA spectrum is steep in the
first epoch, so a spectral turnover would be needed in order to match the
VLASS upper limit, which is not consistent with the host galaxy emission (that
should be steep).

Table 3: ATCA radio observations of AT 2022dsb. MJD |  Date |  Array config. |  Frequency (GHz) |  Flux density ($\mu$Jy)
---|---|---|---|---
59661 |  2022-03-23 |  6A |  5.5 |  $593\pm 19$
59661 |  2022-03-23 |  6A |  9 |  $536\pm 17$
59819 |  2022-08-28 |  6D |  5.5 |  211$\pm$11
59819 |  2022-08-28 |  6D |  9 |  152$\pm$10
59912 |  2022-11-29 |  6C |  5.5 |  171$\pm$8
59912 |  2022-11-29 |  6C |  9 |  127$\pm$6
Figure 7: 5.5 GHz radio luminosity of AT2022dsb (red stars) compared to a
selection of other radio-detected thermal TDEs (AT 2019azh, Goodwin et al.
2022a; AT 2020opy, Goodwin et al. 2022b; AT 2019dsg, Cendes et al. 2021;
ASASSN 14li, Alexander et al. 2016; ASASSN 15oi, Horesh et al. 2021; CNSS
J0019+00, Anderson et al. 2020; XMMSL1 J0740-85, Alexander et al. 2017; IGR
J12580+0134, Irwin et al. 2018; AT2020vwl Goodwin et al. 2023; AT2018hyz
Cendes et al. 2022). The horizontal axis indicates the time since first
detection of the source at optical or X-ray wavelengths.

##  4 Data analysis

###  4.1 X-ray spectral fitting

The X-ray spectra were analysed using the Bayesian X-ray Analysis software
(BXA; Buchner et al. 2014), which connects the nested sampling algorithm
UltraNest101010https://johannesbuchner.github.io/UltraNest/ (Buchner, 2021)
with the fitting environment CIAO/Sherpa (Fruscione et al., 2006). The eROSITA
and XMM PN spectra were fitted in the 0.2–8 keV and 0.2–10 keV range,
respectively. A joint fit of the source and background spectra was performed,
using the C-statistic for fitting (Cash, 1976), and modelling the background
using the principal component analysis (PCA) technique described in Simmonds
et al. (2018). The Galactic absorption is modelled with a total (HI and H2)
Galactic hydrogen column density of $1.73\times 10^{21}$ cm-2 (Willingale et
al., 2013), cosmic abundances from Wilms et al. (2000) and cross sections from
Verner et al. (1996).

Each of the eROSITA and XMM PN spectra were fitted with the following source
models, commonly used to fit the X-ray spectra of TDEs: i) zbbody: redshifted
blackbody ii) zpowerlaw: redshifted power-law, iii) zbremsstrahlung:
redshifted thermal bremsstrahlung. To assess the goodness of fit and compare
between different fitted models, we use the Akaike Information Criterion
($AIC$), defined as $AIC=2k-2\ln{\hat{\mathcal{L}}}$, where $k$ is the number
of free-parameters in the fitted model, and $\hat{\mathcal{L}}$ the estimated
maximum likelihood from the spectrum fitting; the lower the value of the AIC,
the better the fit to the spectrum. An overview of the spectral fit parameters
are listed in Table 4.

The eRASS5 spectrum, obtained $\sim$14 days before the optical peak, is ultra-
soft, and can be well fitted by the thermal bremsstrahlung model with
temperature $kT_{\mathrm{brems}}=71^{+8}_{-5}$ eV (Fig. 8), or a blackbody
with temperature $kT_{\mathrm{bb}}=47^{+5}_{-5}$ eV; such temperatures are
consistent with the X-ray spectra of other X-ray bright thermal TDEs (e.g.
Saxton et al. 2020). This corresponds to a 0.2–2 keV observed flux,
$F_{\mathrm{X,\,obs}}=(3.4^{+0.6}_{-0.5})\times 10^{-13}$ erg cm-2 s-1, and a
0.2–2 keV unabsorbed flux, $F_{\mathrm{X,\,unabs}}=(1.6^{+0.4}_{-0.4})\times
10^{-11}$ erg cm-2 s-1($L_{\mathrm{X}}=$$2.5^{+0.6}_{-0.5}\times 10^{43}$ erg
s -1).

Figure 8: BXA fitted models to the convolved X-ray spectra from eROSITA (top,
14 days before optical peak) and XMM (bottom two plots, $\sim$5 days and
$\sim$180 days after optical peak). Black and grey markers represent source
and scaled background spectra, respectively, with the background component not
originating from the TDE host galaxy. The solid red line denotes the median
model, whilst the shaded red band encloses 68% of the posterior. The preferred
model (Table 4) for the eRASS5 spectrum is thermal bremsstrahlung with
$kT_{\mathrm{brems}}=71^{+8}_{-5}$ eV, whilst it is a powerlaw for the XMM
spectra, with $\Gamma$ of $2.7^{+0.3}_{-0.3}$ and $3.5^{+0.5}_{-0.5}$,
respectively. The unconvolved models and spectra are presented in Fig. 9.

The first XMM PN spectrum, taken $\sim$19 days after the eRASS5 spectrum, is
harder, and can be best-fit by a power-law with photon index
$2.7^{+0.3}_{-0.3}$. The eRASS5 to XMM spectral hardening is also accompanied
by a factor of $\sim$39 decrease in $F_{\mathrm{X,\,obs}}$ to
$(8.9^{+2.3}_{-1.8})\times 10^{-15}$ erg cm-2 s-1. The early-time evolution of
the spectral energy distribution (SED) evolution between the eROSITA and XMM
observation is plotted in Fig. 9. At $\sim$173 ($\sim$154) days after the
eRASS5 (first XMM observation), the second XMM observation shows a power-law
slope consistent with the first XMM observation with photon index
$3.5^{+0.5}_{-0.5}$, as well as a similar observed 0.2–2 keV flux of
$(6.6^{+1.3}_{-1.1})\times 10^{-15}$ erg cm-2 s-1. The 0.2–2 keV fluxes in
these XMM observations are below the 3$\sigma$ upper limit inferred from the
stacked eROSITA observations from its first four all-sky surveys (section 2).

Table 4: X-ray spectral fit results. The Akaike Information Criterion ($AIC$) column estimates the goodness of fit, with a lower $AIC$ representing a better fit. MJD | ObsID | zbbody | zpowerlaw | zbremsstrahlung
---|---|---|---|---
|  | AIC |  $kT$ [eV] | AIC | $\Gamma$ | AIC |  $kT_{\mathrm{brems}}$ [eV]
59627.939 | eRASS5 | 367.6 | $47^{+5}_{-5}$ | 367.9 | $7.7^{+0.2}_{-0.4}$ | 366.8 | $71^{+8}_{-5}$
59646.226 | XMM1 | 9908.3 | $219^{+28}_{-24}$ | 9893.8 | $2.7^{+0.3}_{-0.3}$ | 9899.7 | $768^{+346}_{-173}$
59820.520 | XMM2 | 7925.8 | $150^{+26}_{-23}$ | 7918.0 | $3.5^{+0.5}_{-0.5}$ | 7921.1 | $403^{+155}_{-96}$
Figure 9: Early SED evolution of AT 2022dsb. The red and orange markers show
the SED at the time of the eRASS5 detection (MJD$\sim$59627, 14 days before
optical peak) and first XMM observation $\sim$19 days later (MJD$\sim$59646, 5
days after optical peak). The dotted and solid lines in the X-ray band-pass
(grey region) denote the observed and unabsorbed best fitting spectral models.
The two blackbody curves (blue) passing through the ATLAS $o$-band data point
on MJD$\sim$59627 are at temperatures of $\log[T/\mathrm{K}]=4.1$ and $4.56$,
the minimum and maximum temperatures inferred from fitting a single
temperature blackbody to the multi-band photometry of a ZTF-selected TDE
population (Hammerstein et al., 2023).

The photon index of $\Gamma\sim 2.7$ in the first XMM observation is much
softer than the photon indices of X-ray bright AGN; for example, Nandra &
Pounds (1994) characterised a sample of continuum slopes of Seyfert galaxies
by a Gaussian with $1.95\pm 0.15$. In addition, it is also softer than the
hard X-ray emission from an advection-dominated accretion flow (ADAF; e.g.
Narayan et al. 1998) of slope $\lesssim 2$ (Gu & Cao, 2009), yet harder than
the spectra of thermal TDEs (Saxton et al., 2021). With a 0.2–2 keV luminosity
of $\sim 4\times 10^{40}$ erg s-1 and with no major change in flux in the
0.2–2 keV band between the two XMM epochs, we consider the XMM source spectra
to be likely dominated by diffuse X-ray emission from within the circumnuclear
environment of the host galaxy (i.e. unrelated to the TDE-triggered accretion
episode onto the SMBH). This is in part motivated by the host galaxy of AT
2022dsb likely hosting a LLAGN prior to its 2022 outburst (recall the mid-
infrared colour of $W1-W2\sim 0$ and previous type II AGN classification for
its host galaxy; section 2), with past X-ray observations of nearby LLAGN also
suggesting the presence of hot, diffuse plasma within a few hundred parsecs of
the nucleus (Flohic et al., 2006), and with $\log[L_{0.5-2\,\mathrm{keV}}]$
and $\log[L_{2-10\,\mathrm{keV}}]$ luminosities in the range of $40.2\pm 1.3$
and $39.9\pm 1.3$ (González-Martín et al., 2009).

A similar scenario may also have been present in the TDE candidate ASASSN-15oi
(Gezari et al., 2017), where two XMM spectra, taken $\sim$80 days and
$\sim$230 days after optical discovery, were best fitted111111These spectral
fit results are reported in Gezari et al. 2017 and were obtained using the
first XMM spectrum. by a two component model consisting of a blackbody with
$kT=47\pm 3$ eV and a power-law with $\Gamma=2.5\pm 0.8$. As ASASSN-15oi
brightened in the X-rays over the $\sim 160$ days between these spectra, only
the normalisation of the blackbody component increased, without a significant
change in $kT$ or $\Gamma$. This would require a fair amount of model fine-
tuning if the power-law component originated from Compton upscattered TDE disc
photons, and the disc luminosity varied over time. Instead, this may be more
easily explained if the disc emission evolves independently of the lower
luminosity diffuse host emission (as also suggested in Gezari et al. 2017),
with the latter being emitted at much larger physical length scales than the
X-ray emission from the TDE disc.

If the XMM source spectra are dominated by the diffuse emission with the host
galaxy, then the soft X-ray emission dominating the eRASS5 spectrum may have
been obscured by optically thick material (see section 6 for further
discussion on its possible origin). Assuming that the TDE disc has spectral
properties in its first XMM spectrum similar to eRASS5 ($kT_{\mathrm{BB}}\sim
50$ eV) and a similar blackbody normalisation, then an increase in the neutral
hydrogen column density along our line-of-sight to $>4.9\times 10^{21}$ cm-2
would be capable of causing a 0.2—2 keV flux drop by a factor of at least 39
(between these two spectra), or a He ii column density of $1.6\times 10^{21}$
cm-2 if the flux drop is due to photoionisation of He ii (Fig. 10), when
modelling He ii absorption121212The column densities of all other species are
set to zero within this absorption model here. using the xspec model ISMabs
(Gatuzz et al., 2015). The conversion of these column densities into an
estimate of the mass of a debris envelope obscuring the disc is complicated by
the unknown ionisation fraction of helium in the debris; an alternate approach
to constrain the reprocessor’s properties at early times is presented in
section 6.

Figure 10: The amplitude of the observed flux drop in the 0.2–2 keV band
($F_{\mathrm{X,unabs}}/F_{\mathrm{X,obs}}$) due to absorption by neutral
hydrogen (top) and He ii (bottom), with each curve representing a different
blackbody temperature ($kT\sim 0.05$ keV for AT 2022dsb). The sensitivity of
$F_{\mathrm{X,unabs}}/F_{\mathrm{X,obs}}$ to absorption depending on the
balance between the ionisation potential of the absorber and $kT$. The
horizontal red dashed line corresponds to the flux drop by a factor of 39 at
early times seen in AT 2022dsb, corresponding to an $N_{\mathrm{H}}$
($N_{\mathrm{HeII}}$) of $4.9\times 10^{21}$ cm-2 ($1.6\times 10^{21}$ cm-2).

###  4.2 Photometric analysis

Following Malyali et al. (2023b), the ATLAS $o$-band lightcurves were fitted
with a half-Gaussian rise, exponential decay model, as described in van Velzen
et al. (2019), and plotted in Fig. 11. Only the $o$-band photometry was fitted
here due to this providing the best sampling of the rise and decay of the
optical emission, and only photometry in the range 59600$<$MJD$<$59750 was
used. The inferred rise and decay timescales are $\sigma=7.9^{+0.4}_{-0.4}$
days and $\tau=21.0^{+0.6}_{-0.6}$ days, respectively, with the lightcurve
peaking at MJD=$59640.9^{+0.5}_{-0.4}$; this value is used as the reference
time for the optical peak of AT 2022dsb in this work. The peak inferred
$F_{\nu}$ is $320^{+7}_{-6}$ $\mu$Jy, corresponding to $\nu L_{\nu}\sim
2\times 10^{42}$ erg s-1. To help understand when the eRASS5 detection occurs
during the evolution of the TDE, then it is also valuable to define a start
time for the optical rise based on the fitted model here. If one considers
this to be when the optical flux is $\sim$1% ($\sim 0.03$%) of the optical
peak, then this would occur when $\mathrm{MJD}_{\mathrm{start}}\sim$59617
($\sim 59609$).

Figure 11: Half-Gaussian rise, exponential decay model (red) fitted to the
ATLAS $o$-band photometry (orange markers). The shaded red bands denote the
credible region enclosed by the 16th and 84th percentiles of the posterior.
The datapoint with $F\sim 300$ $\mu$Jy at MJD$\sim$59600 has a 1$\sigma$ error
consistent with 0 (i.e. is not precursor emission to the main flare).

To obtain a more physically-motivated estimate of the start time of the event,
then we also fitted the multi-band photometry (ATLAS $o$ and $c$, ZTF $g$ and
$r$, Swift $UVW1$, $UVM2$ and $UVW2$ bands; Fig. 12) with the TDE module
(Mockler et al., 2019) of the Modular OpenSource Fitter for Transients
(MOSFiT; Guillochon et al. 2018), using the nested sampler dynesty for
posterior sampling (Speagle, 2020). The free parameters of this model are the
black hole mass ($M_{\mathrm{BH}}$), mass of the disrupted star ($M_{\star}$),
the scaled impact parameter ($b$), the efficiency of converting accretion
luminosity into the optical luminosity ($\epsilon$), a normalising factor and
exponent for the photosphere radius ($R_{\mathrm{ph,0}}$, $l_{\mathrm{ph}}$),
a viscous delay timescale ($T_{\mathrm{viscous}}$) and the time of first
massfall back to pericentre ($\mathrm{MJD}_{\mathrm{fb}}$). Although most of
the inferred parameter estimates are dominated by systematics (Table 5), we
note that the MOSFiT modelling does suggest a lower mass SMBH for the
disruption with $\log(M_{\mathrm{bh}}/M_{\odot})=(6.4^{+0.2}_{-0.2})\pm$0.2,
as compared to the black hole mass derived from the
$M_{\mathrm{BH}}-M_{\star}$ relation (Reines & Volonteri, 2015) using
$M_{\star}$ from the host galaxy SED analysis (section 2). The inferred
$\mathrm{MJD}_{\mathrm{fb}}=(59610.2^{+6.6}_{-6.3})\pm 15$ is consistent with
the start time for the optical rise inferred above.

Figure 12: MOSFiT model fits to the multi-band photometry of AT 2022dsb, with colour scheme following Fig. 3. The black and red lines mark the estimated median time of first mass fallback ($\mathrm{MJD}_{\mathrm{fb}}=59610\pm 6$) and the eRASS5 coverage of AT 2022dsb, occurring only $\sim$17 days later. A zoom-in on the ATLAS o-band difference photometry sampling the optical rise, and which is used for constraining $\mathrm{MJD}_{\mathrm{fb}}$, is presented in Fig. 11. Table 5: Posterior medians and 1$\sigma$ credible regions inferred from the MOSFiT TDE lightcurve fitting. The estimated systematic errors on each estimate are taken from Mockler et al. (2019). Parameter | Value | Systematic Error
---|---|---
$\log(M_{\mathrm{bh}}/M_{\odot})$ | $6.4^{+0.2}_{-0.2}$ |  $\pm$0.2
$\log(M_{\star}/M_{\odot})$ | $0.2^{+0.3}_{-0.1}$ |  $\pm$0.66
$b$ | $1.0^{+0.2}_{-0.3}$ |  $\pm$0.35
$\log(\epsilon)$ | $-2.0^{+0.5}_{-0.5}$ |  $\pm$0.68
$\log(R_{\mathrm{ph,0}})$ | $-0.1^{+0.3}_{-0.3}$ |  $\pm$0.4
$l_{\mathrm{ph}}$ | $1.0^{+0.2}_{-0.2}$ |  $\pm$0.2
$\log(T_{\mathrm{viscous}}/\mathrm{days})$ | $-0.9^{+1.1}_{-1.2}$ |  $\pm$0.1
$\mathrm{MJD}_{\mathrm{fb}}$ | $59610.2^{+6.6}_{-6.3}$ |  $\pm$15

###  4.3 Spectroscopic analysis

We used a modified version of the Python quasar fitting code (PyQSOFit; Guo et
al. 2018) to fit the optical spectra of AT 2022dsb after de-reddening the
Galactic foreground contribution. First, we fitted the emission line free
regions of each of the optical spectra with a power-law to estimate the
continuum contribution to the spectrum, before dividing the observed spectrum
by this continuum component to obtain a normalised spectrum. We note that we
do not attempt to model or subtract the host galaxy component here, such that
the modelled continuum involves both TDE and host galaxy emission. All
spectral fitting made use of the python package lmfit (Newville et al., 2014)
and the Markov Chain Monte Carlo (MCMC) sampler emcee (Foreman-Mackey et al.,
2013).

After normalising by the continuum, we fitted each of the narrow emission
lines in this complex (H$\alpha$, [N ii] 6548Å and 6583 Å) with a single
Gaussian, and forced each of these to be of the same width. The broad
H$\alpha$ component was fit with a single Gaussian (Fig 13). Due to the
possible presence and blending of emission from H$\beta$, He ii 4686Å, N iii
4640Å, H$\gamma$, Fe ii 37, 38, within the broad He ii complex, it is not
straightforward to constrain the evolution of each of these possible
components; we therefore examine the more isolated broad H$\alpha$ emission
lines here.

Figure 13: Example fit to the H$\alpha$ complex, with the continuum normalised
spectrum in black and the best fitting model in red. The grey band denotes a
region of telluric absorption which was masked during fitting. The centroid of
the broad H$\alpha$ is shifted by $-1600\pm 300$ km s-1 with respect to the
rest frame of the host galaxy.

In the first LCO FLOYDS spectrum obtained $\sim$5 days before optical peak,
the full-width half max (FWHM) of the broad H$\alpha$ line is $10400\pm 400$
km s-1, and its centroid is blueshifted to $\lambda_{\mathrm{rest}}=$$6530\pm
7$ Å, corresponding to a velocity of $-1600\pm 300$ km s-1. In the second
optical spectrum obtained $\sim$1 day before optical peak, the FWHM is
$10500\pm 300$ km s-1, consistent with the earlier spectrum, but the velocity
offset is $200\pm 400$ km s-1. At later times, the H$\alpha$ emission line is
not clearly seen above the host galaxy continuum emission.

##  5 early-time X-ray emission in TDEs

In the following section, we briefly review the literature on early-time X-ray
observations of TDEs, and compare the X-ray transient seen in AT 2022dsb (Fig.
3) with the wider TDE population.

The majority of the X-ray selected TDE population known prior to the launch of
eROSITA was first discovered at a time when wide-field, high-cadence optical
surveys were still relatively limited compared with the current generation,
with respect to their depth, cadence, sky coverage, difference-imaging
capabilities (important for nuclear transients), and ease-of-access to their
optical lightcurves. Largely as a result of this, only a handful of the X-ray
selected TDE candidates showed transient optical/ UV emission (Saxton et al.,
2020), which was only ever identified after the initial X-ray discovery, and
typically only through Swift UVOT follow-up. Furthermore, of the systems with
detected transient optical/ UV emission, only the decaying phase of the TDE
lightcurve was sampled in the UV (see discussion in Section 9 in a recent
review Saxton et al. 2020), thus the early-time X-ray evolution (during the
initial optical/ UV brightening) of these X-ray selected systems remains
unknown. Whilst the launch of eROSITA has seen a vast increase in the number
of X-ray bright TDE candidates (e.g. Malyali et al. 2021, 2023a, 2023b; Liu et
al. 2023; Homan et al. 2023), the majority of the TDE candidate population
show no transient optical emission. In a sample of eROSITA-selected TDE
candidates (Sazonov et al., 2021), only four systems display both transient
X-ray and optical emission, but the detections of flaring X-ray emission
associated to the TDE always occurs after the optical peak.

Although the number of optically-selected TDE candidates has rapidly increased
over the last decade, the majority of these have X-ray observations commencing
at earliest near to, or after, peak optical brightness. This may in part stem
from the very high discovery rate of transients in the latest generation of
optical surveys, meaning that astronomers generally wait until close to peak
optical brightness for an optical transient to become a strong TDE candidate,
and only then trigger X-ray follow-up observations. Despite this, there are
still six optically-bright TDEs with X-ray observations starting before
optical maximum (Table 6 and Fig. 14), where this list was obtained via visual
inspection of the joint X-ray and optical lightcurves of the TDEs presented in
both the ZTF TDE sample (Hammerstein et al., 2023), and those in the recent
TDE review paper by Gezari (2021). All of these TDEs are of H+He
type131313This is likely due current TDE follow-up strategies, since if a He
complex is detected in a follow-up optical spectrum obtained during the
optical rise, then there’s a stronger indication that the event is a TDE at
these early times, and a higher likelihood of Swift XRT and UVOT observations
being triggered with high urgency to monitor the evolution of the system. ,
with the exception of AT 2019ahk, which only shows transient broad Balmer
emission lines. For each of these systems, the 0.3–2 keV XRT lightcurves were
generated and downloaded from the UKSSDC as in section 3.1.3. These were then
converted to 0.2–2 keV lightcurves using webPIMMS, assuming a redshifted
blackbody spectrum with $kT=50$ eV (similar to other X-ray bright TDEs; Saxton
et al. 2020), absorbed by a Galactic $N_{\mathrm{H}}$ along the line-of-sight
to the TDE taken from Willingale et al. (2013); lower $kT$ values for each TDE
here would lead to higher estimated 0.2–2 keV fluxes.

Table 6: Properties of TDEs with X-ray observations pre-optical peak. MJDpeak is the inferred optical peak, $L^{\mathrm{1st}}_{\mathrm{X}}$ is the inferred 3$\sigma$ upper limit on the 0.2–2 keV luminosity from the earliest X-ray observation (by Swift XRT) and $P$ denotes the phase of the first X-ray observation relative to MJDpeak. $\sigma$ is the inferred rise timescale from a half-Gaussian fitted to the optical lightcurve (section 4.2). MJDpeak and $\sigma$ were taken from van Velzen et al. (2020) for iPTF15af, AT 2018dyb, AT 2019ahk, AT 2019azh, AT 2019qiz, and Hammerstein et al. (2023) for AT 2020zso. Name | MJDpeak |  $L^{\mathrm{1st}}_{\mathrm{X}}$ [erg s-1] |  $P$ [d] |  $\sigma$ [d] | $P/\sigma$
---|---|---|---|---|---
AT2018dyb | $58340.7^{+1.4}_{-1.3}$ | $<7\times 10^{42}$ | $-23.1$ | $31.6^{+2.3}_{-0.7}$ | $-0.7$
AT2019ahk | $58548.3^{+0.8}_{-0.9}$ | $<3\times 10^{42}$ | $-33.5$ | $20.0^{+0.5}_{-0.5}$ | $-1.7$
iPTF15af | $57061.0^{+1.6}_{-1.8}$ | $<2\times 10^{43}$ | $-11.7$ | $31.6^{+1.5}_{-2.1}$ | $-0.4$
AT2019azh | $58558.6^{+1.5}_{-1.6}$ | $<1\times 10^{42}$ | $-13.8$ | $20.0^{+2.4}_{-1.3}$ | $-0.7$
AT2019qiz | $58761.4^{+0.6}_{-0.6}$ | $<1\times 10^{42}$ | $-8.3$ | $6.3^{+0.1}_{-0.1}$ | $-1.3$
AT2020zso | $59188.0^{+1.4}_{-1.4}$ | $<4\times 10^{43}$ | $-15.0$ | $6.9^{+0.2}_{-0.2}$ | $-2.2$
AT2022dsb | $59640.9^{+0.5}_{-0.4}$ | $2.5^{+0.6}_{-0.5}\times 10^{43}$ | $-13.5$ | $7.9^{+0.4}_{-0.4}$ | $-1.7$
Table 7: Reported outflow properties of TDEs with X-ray observations pre-optical peak. Each of these objects have been followed up with different instruments and at different times, so there is not necessarily a common observed outflow indicator between systems. The expanding photosphere at early times points to a mechanical outflow only if it traces the motion of the debris (e.g. Nicholl et al. 2020). Name |  Outflow properties |  Reference
---|---|---
AT2018dyb |  H$\alpha$ blueshifted by $\sim$700 km s-1 in spectrum obtained $\sim$24 days before peak (although H$\alpha$ redshifted after optical peak). |  Leloudas et al. (2019)
AT2019ahk |  Photosphere expanding at $\sim$2700 km s-1 via modelling photometric rise, assuming constant temperature. |  Holoien et al. (2019)
iPTF15af |  No H$\alpha$ detected in first optical spectrum near peak. Si iv blueshifted by $\sim$6000 km s-1 in HST spectrum obtained $\sim$28 days post-peak. |  Blagorodnova et al. (2019)
AT2019azh |  Transient radio emission detected $\sim$10 days before optical peak. |  Goodwin et al. (2022a)
AT2019qiz |  Photosphere expanding at $\sim$2200 km s-1 via modelling photometric rise, assuming constant temperature. Inferred outflow velocity $\lesssim$5000 km s-1 from asymmetric H$\alpha$ line in spectrum $\sim$9 days before peak. |  Nicholl et al. (2020)
AT2020zso |  Photosphere expanding at $\sim$2900 km s-1 via modelling photometric rise, assuming constant temperature. |  Wevers et al. (2022)

Figure 14: Early-time X-ray lightcurves of TDEs with observations pre-optical
peak, with markers following the same definition as for Fig. 3. The dotted red
lines mark the phase of the eRASS5 observation of AT 2022dsb and inferred
0.2–2 keV luminosity, whilst the dark orange solid line marks the eRASS5
observation at the time of its normed phase, $P/\sigma$ (Table 6). MJDs are
defined with respect to the inferred optical peaks in Table 6. The eROSITA
observation clearly represents the earliest X-ray detection of a TDE to date,
and although there are XRT observations of these other TDEs at a comparable
phase, only AT 2020zso has been observed at an earlier $P/\sigma$.

Comparing the X-ray lightcurve of AT 2022dsb with the other TDEs with pre-peak
X-ray observations (Fig. 14), then it is clear that the early-time transient
X-ray emission in AT 2022dsb has never been observed before across all known
TDEs with well sampled optical peaks. Of the TDEs in Table 6, only AT 2019azh
has a significant detection of soft X-ray emission before the observed optical
maximum, with the system being detected for the first time $\sim$3 days before
optical peak. AT 2019azh then remains at approximately a constant
$L_{\mathrm{X}}$ over the following $\sim$40 days after the first significant
detection (Fig. 14), and thus shows a vastly different X-ray evolution to AT
2022dsb.

The eRASS5 observation is not the earliest X-ray observation of a TDE in terms
of the phase (number of days observed before optical maximum), with AT
2018dyb, AT 2019ahk, AT 2019azh and AT 2020zso all having been observed at
earlier phases than AT 2022dsb. Furthermore, each of the earliest time
observations for each system should also have been able to detect AT 2022dsb-
like X-ray emission, given the 3$\sigma$ upper limits on $L_{\mathrm{X}}$ were
lower than the $L_{\mathrm{X}}$ of the eRASS5 observation of AT 2022dsb.
However, the optical-UV lightcurves of these TDEs evolve differently to AT
2022dsb, with respect to their peak luminosities and rise timescales. This is
not unexpected, since these systems may span a range of different black hole
masses, stellar masses and impact parameters. For example, the rise timescale
of AT 2022dsb is inferred to be $\sim$8 days, whereas it is $\sim$31 days for
iPTF-15af (van Velzen et al., 2020) (Table 6). This complicates a clean
comparison between these systems, and the task of understanding how early on
in a TDE’s evolution an AT 2022dsb-like X-ray transient may be observable. If
one considers the normalised phase (phase divided by the estimated rise
timescale; Table 6), then the eRASS5 observation of AT 2022dsb represents the
second earliest X-ray observation of a TDE showing a transient He ii emission
complex, with only AT 2020zso being observed at an earlier stage of the
lightcurve.

Whilst the detection of the early-time X-ray transient in AT 2022dsb certainly
benefitted from eROSITA serendipitously scanning over it a day before the
first optical detection, this cannot be the sole factor in this discovery
given the XRT coverage of other TDEs described above (i.e. there were
observations that were early and deep enough to detect a source with spectral
properties similar to AT 2022dsb, but did not because of physical differences
between these systems). The fact that the X-ray emission could have been
observed in other TDEs but was not, particularly for AT 2020zso (the TDE with
observations performed at the earliest normed phase), suggests that the
assumptions made when converting the observed 0.3–2 keV XRT count rate into an
unabsorbed 0.2–2 keV flux (corrected only for Galactic absorption), and then a
0.2–2 keV intrinsic luminosity, may have been oversimplified and require
further consideration. For example, an additional absorber along the line-of-
sight to the TDE disc (such as from stellar debris ejected during the
circularisation process, which may be optically thick or thin depending on the
observer’s viewing angle to the system, or from neutral hydrogen in the host
galaxy unrelated to the TDE), or a lower effective temperature of the disc
emission (e.g. due to a retrograde black hole spin), would lead to larger
estimated unabsorbed flux upper limits from the XRT observations, and might
explain the previous non-detections of X-ray emission in these systems. The
early-time X-ray transient seen in AT 2022dsb is likely not a universally
observable feature in TDEs.

##  6 Discussion

The unique observational feature of AT 2022dsb with respect to the wider
population of optically-selected TDEs is its early-time transient X-ray
emission detected by eROSITA (see Table 8 for a summary of the key events in
the evolution of AT 2022dsb). From the physical modelling of the multi-band
photometry (section 4.2), then the time of first mass fallback to pericentre
after the disruption is estimated to be MJD$\sim 59610$ ($\sim$31 days before
optical peak), meaning that the eROSITA discovery of ultra-soft X-ray emission
on MJD 59627 occurs only 17 days after this. As the optical emission brightens
in the system, then the observed X-ray emission in the 0.2–2 keV band decays
over a 19 day period by a factor of $\sim$39 (Fig. 3). This joint X-ray-to-
optical evolution has not been observed before in a TDE candidate. Although
the observed X-ray emission rapidly dims during the optical rise, AT 2022dsb
shows a broad He ii complex which persists for at least $\sim$38 days after
optical peak, and was first detected 5 days before peak (Fig. 6). Several
outflow signatures are also present during the early stages of this TDE, in
the form of blueshifted H$\alpha$ at $-1600\pm 300$ km s-1, first observed at
$P=-5$ days (section 4.1), radio transient emission at $P=20$ days (Fig. 7),
and blueshifted Ly$\alpha$ absorption lines at $\sim$3000 km s-1, observed at
$P=54$ days (Engelthaler & Maksym 2023, Engelthaler et al., in prep.).
Importantly, other past X-ray observations of He-TDEs before optical peak have
not detected X-ray emission at a similar $L_{\mathrm{X}}$ (Table 6), despite
the observations being carried out at a similar phase to the eRASS5
observation of AT 2022dsb, and also having upper limits on $L_{\mathrm{X}}$
lower than for the eRASS5 detection of AT 2022dsb; each of the He TDEs in this
sample have also been reported to show outflow signatures in observations
performed around optical peak (Table 7).

MJD |  Event
---|---
$\sim$59610 |  Time of first mass fallback to pericentre (section 4.2).
59627 |  eRASS5 detection.
59628 |  First 3$\sigma$ detection of optical emission in ATLAS $o$-band.
59636 |  Broad blueshifted H$\alpha$ and an emission complex detected around He ii 4686Å, in the first follow-up optical spectrum of AT 2022dsb.
59641 |  Peak in the observed optical flux (section 4.2).
59646 |  First XMM observation, 0.2–2 keV observed flux drop by a factor of $\sim$39 relative to eRASS5.
59661 |  Detection of radio transient emission with ATCA in first radio follow-up observation.
59693 |  First detection of outflow at 3000 km s-1 from Ly$\alpha$ absorption (FWHM$\sim 14000$ km s-1) in first HST spectrum (Engelthaler & Maksym, 2023).
Table 8: Key events in the early evolution of AT 2022dsb.

The early X-ray emission detected by eROSITA likely comes from an accretion
disc that has recently been assembled through circularisation of the earliest-
arriving gas in the fallback stream (Bonnerot et al., 2021). We rule out the
early X-ray transient being produced by shock breakout emission from the
surface of the star after being maximally compressed at pericentre (Carter &
Luminet, 1983; Guillochon et al., 2009; Stone et al., 2013; Yalinewich et al.,
2019), as the predicted timescales of $\mathcal{O}$(minutes) for these flares
are far shorter than what is observed in the eRASS5 observation ($>$1 day;
Fig. 5). We also disfavour the X-ray transient being caused by accretion disc
cooling (e.g. Cannizzaro et al. 2021), since the disc temperature should
increase as the optical emission brightens if the optical light curve traces
the accretion rate at early times, or Lense-Thirring driven precession of the
newly formed disc (Stone & Loeb, 2012; Franchini et al., 2016), as no
rebrightening episodes are detected over the X-ray follow-up campaign (Fig.
3). Regardless of origin, then the high-energy tail of this early hard
ionising source also likely drives the ionisation of He ii and the formation
of the He-complex in the optical spectra.

Around 14 days before the optical peak, it is likely that we have an
unobscured view onto the nascent disc, which is initially surrounded by an
envelope of gas formed through shocks during the circularization process, as
found in simulations (Bonnerot et al., 2021). While this envelope is initially
of low enough density for the disc emission to promptly emerge, the envelope
mass increases over time due to feeding by the outflowing gas in the system.
As a result, the X-ray emission may be more efficiently absorbed over time,
leading to the observed drop-off in the X-ray emission, and a reprocessed-
driven optical brightening. To quantitatively test this scenario, we follow
Roth et al. (2016) by modelling the envelope as a sphere of inner and outer
radii $R_{\rm in}$ and $R_{\rm out}$, containing a mass of gas $M_{\rm env}$
distributed according to a density profile $\rho\propto R^{-2}$, which is
irradiated by an inner source of luminosity $L$. The effective optical depth
$\tau_{\rm eff}$ relevant for X-ray absorption is then given by their equation
27, relying on He ii photoionization being the dominant absorption process and
using solar composition. Here, we further assume that the envelope mass
increases with time as $M_{\rm env}=\dot{M}_{\rm env}t$ due to feeding by
early returning debris. The time at which the envelope is able to absorb the
inner X-ray radiation is obtained by solving $\tau_{\rm eff}(t)=1$, which
gives

$t_{\rm abs}=24\,{\rm
d}\,\left(\frac{L}{10^{43}\,\rm{erg\,s^{-1}}}\right)^{5/19}\left(\frac{\dot{M}_{\rm
env}}{M_{\odot}\,\rm{yr^{-1}}}\right)^{-1}\\\ \left(\frac{R_{\rm
in}}{10^{14}\,\rm{cm}}\right)^{20/19}\left(\frac{R_{\rm
out}}{10^{15}\,\rm{cm}}\right)^{16/19},$ (1)

where the luminosity $L\approx 10^{43}\,\rm{erg\,s^{-1}}$ is determined from
that of the detected X-rays. This time is approximately consistent with the 19
day delay between the eRASS5 detection and first XMM observation, for an
envelope feeding rate $\dot{M}_{\rm env}\approx M_{\odot}\,\rm{yr^{-1}}$,
comparable to the debris fallback rate of a typical TDE near peak (e.g. Rossi
et al. 2021). Within this interpretation, the early-time X-ray detection
presented in this paper provides a new way to constrain physical properties
such as the feeding rate and size of the envelope, and the luminosity of the
obscured accretion disc, which are crucial to improve our theoretical
understanding of these systems.

The origin of the early-time outflows seen in some TDEs is still unclear (see
discussion in Goodwin et al. 2022a), but is thought to be due to either a
stream-stream collision-induced outflow (CIO; e.g. Lu & Bonnerot 2020), debris
unbounded by accretion luminosity (Metzger & Stone, 2016) or a radiatively-
driven disc wind (Lodato & Rossi, 2011; Miller, 2015). The data set in this
work does not allow us to distinguish between these mechanisms for AT 2022dsb,
since each scenario would initially lead to an increased density of gas and
optical depth along our line-of-sight as the fallback rate increases over
time141414Future spectroscopic monitoring of TDEs in the UV may distinguish
between these two origins..

Importantly, if the observed early-time evolution in AT 2022dsb is not unique
to this system, then other TDEs may also be X-ray bright at early times, and
become X-ray faint only when veiled by outflowing debris launched shortly
after the onset of circularisation.151515Alternatively, if the X-ray emission
is above a critical luminosity, then all of the He ii may be ionised to He
iii, enabling the X-ray emission to escape the system (Metzger & Stone,
2016).. Given that the existing models described above predict the launching
of outflows which would extend large solid angles on the sky, as seen by the
disrupting black hole, then a large fraction of optically-bright TDEs may
therefore be X-ray faint when followed up in the weeks-to-months after optical
peak (unless viewed at angles peering through an optically thin funnel in the
reprocessor; Metzger & Stone 2016; Dai et al. 2018; Lu & Bonnerot 2020).

##  7 Summary

We reported on multi-wavelength observations of the TDE candidate AT 2022dsb,
whose main properties can be summarised as follows:

  1. 1.

eROSITA detected ultra-soft ($kT_{\mathrm{BB}}\sim 45$ eV) X-ray emission
(0.2–2 keV $L_{\mathrm{X}}\sim 3\times 10^{43}$ erg s-1) from a TDE $\sim$14
days before optical peak. The eROSITA detection precedes the first 3$\sigma$
detection in the optical, and occurs only $\sim$17 days after the inferred
time of first mass fallback to pericentre.

  2. 2.

An XMM follow-up observation 19 days after this eROSITA detection revealed a
drop in the observed 0.2–2 keV flux by a factor of 39; during this period, the
optical emission brightened to a maximum. A second XMM observation $\sim$173
days after the eROSITA detection showed a 0.2–2 keV flux and spectral
properties consistent with this first XMM observation. No further X-ray
emission was significantly detected above background by Swift XRT monitoring
observations in the following $\sim$200 days after the eROSITA detection. Thus
without the early-time eROSITA observation, AT 2022dsb would likely have been
classified as an ‘optically bright, X-ray quiet’ TDE.

  3. 3.

Follow-up optical spectra show a broad emission complex around the He ii
4686Å, broad H$\alpha$ emission and a strong blue continuum in the early-time
spectra. The He ii complex is clearly present in the spectra taken $\sim$5
days before optical peak, and is still detected $\sim$38 days after optical
peak (even after the large amplitude X-ray dimming). The strength of these
features with respect to the host galaxy emission decreases over the
spectroscopic follow-up campaign.

  4. 4.

Multiple outflow signatures are detected in the system at early times
(transient radio emission with ATCA, first detected $\sim$20 days post-optical
peak; blueshifted broad H$\alpha$ emission at $\sim$1600 km s-1, detected
$\sim$5 days before optical peak, and blueshifted broad Ly$\alpha$ absorption
at $\sim-$3000 km s-1, detected $\sim$54 days after optical peak).

  5. 5.

The combination of these observed features suggests that outflows launched at
early times may boost the density of the material enshrouding the nascent
disc, leading to an increased amount of reprocessing of the high-energy disc
emission. This causes an early drop-off in the observed X-ray flux whilst the
optical brightens.

  6. 6.

If the observed early-time properties are not unique to this system, then
other TDEs may be X-ray bright at early times, and become X-ray faint when
veiled by outflowing stellar debris. The X-ray vs optically bright nature of a
TDE is also time dependent at early times.

The early-time X-ray emission from TDEs may be monitored in greater detail
with the next-generation of time-domain missions, such as the Einstein Probe
(EP; Yuan et al. 2018), scheduled for launch in late 2023, or through early
follow-up of candidates identified with the Ultraviolet Transient Astronomy
Satellite (ULTRASAT; Sagiv et al. 2014) and the Vera Rubin Observatory
(Ivezic, 2019). High cadence X-ray monitoring observations of such early X-ray
transients may provide a new way to constrain the mass feeding rate and nature
of the reprocessing envelope in TDEs in future work.

## Acknowledgements

AM is grateful to the generosity of Curtin University for hosting his visit,
where parts of this work were completed. AM thanks the XMM, Swift and NICER
teams for approving the ToO requests. AM acknowledges support by DLR under the
grant 50 QR 2110 (XMM_NuTra, PI: Z. Liu). This work was performed in part at
the Aspen Center for Physics, which is supported by National Science
Foundation grant PHY-2210452.

This work is based on data from eROSITA, the soft X-ray instrument aboard SRG,
a joint Russian-German science mission supported by the Russian Space Agency
(Roskosmos), in the interests of the Russian Academy of Sciences represented
by its Space Research Institute (IKI), and the Deutsches Zentrum für Luft- und
Raumfahrt (DLR). The SRG spacecraft was built by Lavochkin Association (NPOL)
and its subcontractors, and is operated by NPOL with support from the Max
Planck Institute for Extraterrestrial Physics (MPE).

The development and construction of the eROSITA X-ray instrument was led by
MPE, with contributions from the Dr. Karl Remeis Observatory Bamberg & ECAP
(FAU Erlangen-Nuernberg), the University of Hamburg Observatory, the Leibniz
Institute for Astrophysics Potsdam (AIP), and the Institute for Astronomy and
Astrophysics of the University of Tübingen, with the support of DLR and the
Max Planck Society. The Argelander Institute for Astronomy of the University
of Bonn and the Ludwig Maximilians Universität Munich also participated in the
science preparation for eROSITA.

The eROSITA data shown here were processed using the eSASS software system
developed by the German eROSITA consortium.

The authors acknowledge support for obtaining the LCO/FLOYDS spectroscopy by
the Deutsche Forschungsgemeinschaft (DFG, German Research Foundation) under
Germany’s Excellence Strategy - EXC-2094 - 390783311.

This work made use of data supplied by the UK Swift Science Data Centre at the
University of Leicester.

The Australia Telescope Compact Array is part of the Australia Telescope
National Facility161616https://ror.org/05qajvd42 which is funded by the
Australian Government for operation as a National Facility managed by CSIRO.
We acknowledge the Gomeroi people as the Traditional Owners of the Observatory
site.

The Legacy Surveys consist of three individual and complementary projects: the
Dark Energy Camera Legacy Survey (DECaLS; Proposal ID 2014B-0404; PIs: David
Schlegel and Arjun Dey), the Beijing-Arizona Sky Survey (BASS; NOAO Prop. ID
2015A-0801; PIs: Zhou Xu and Xiaohui Fan), and the Mayall z-band Legacy Survey
(MzLS; Prop. ID 2016A-0453; PI: Arjun Dey). DECaLS, BASS and MzLS together
include data obtained, respectively, at the Blanco telescope, Cerro Tololo
Inter-American Observatory, NSF’s NOIRLab; the Bok telescope, Steward
Observatory, University of Arizona; and the Mayall telescope, Kitt Peak
National Observatory, NOIRLab. Pipeline processing and analyses of the data
were supported by NOIRLab and the Lawrence Berkeley National Laboratory
(LBNL). The Legacy Surveys project is honored to be permitted to conduct
astronomical research on Iolkam Du’ag (Kitt Peak), a mountain with particular
significance to the Tohono O’odham Nation.

Some of the observations reported in this paper were obtained with the
Southern African Large Telescope (SALT) under the programme 2021-2-LSP-001
(PI: DAHB). Polish participation in SALT is funded by grant No. MEiN
nr2021/WK/01.

NOIRLab is operated by the Association of Universities for Research in
Astronomy (AURA) under a cooperative agreement with the National Science
Foundation. LBNL is managed by the Regents of the University of California
under contract to the U.S. Department of Energy.

This project used data obtained with the Dark Energy Camera (DECam), which was
constructed by the Dark Energy Survey (DES) collaboration. Funding for the DES
Projects has been provided by the U.S. Department of Energy, the U.S. National
Science Foundation, the Ministry of Science and Education of Spain, the
Science and Technology Facilities Council of the United Kingdom, the Higher
Education Funding Council for England, the National Center for Supercomputing
Applications at the University of Illinois at Urbana-Champaign, the Kavli
Institute of Cosmological Physics at the University of Chicago, Center for
Cosmology and Astro-Particle Physics at the Ohio State University, the
Mitchell Institute for Fundamental Physics and Astronomy at Texas A&M
University, Financiadora de Estudos e Projetos, Fundacao Carlos Chagas Filho
de Amparo, Financiadora de Estudos e Projetos, Fundacao Carlos Chagas Filho de
Amparo a Pesquisa do Estado do Rio de Janeiro, Conselho Nacional de
Desenvolvimento Cientifico e Tecnologico and the Ministerio da Ciencia,
Tecnologia e Inovacao, the Deutsche Forschungsgemeinschaft and the
Collaborating Institutions in the Dark Energy Survey. The Collaborating
Institutions are Argonne National Laboratory, the University of California at
Santa Cruz, the University of Cambridge, Centro de Investigaciones
Energeticas, Medioambientales y Tecnologicas-Madrid, the University of
Chicago, University College London, the DES-Brazil Consortium, the University
of Edinburgh, the Eidgenossische Technische Hochschule (ETH) Zurich, Fermi
National Accelerator Laboratory, the University of Illinois at Urbana-
Champaign, the Institut de Ciencies de l’Espai (IEEC/CSIC), the Institut de
Fisica d’Altes Energies, Lawrence Berkeley National Laboratory, the Ludwig
Maximilians Universitat Munchen and the associated Excellence Cluster
Universe, the University of Michigan, NSF’s NOIRLab, the University of
Nottingham, the Ohio State University, the University of Pennsylvania, the
University of Portsmouth, SLAC National Accelerator Laboratory, Stanford
University, the University of Sussex, and Texas A&M University.

BASS is a key project of the Telescope Access Program (TAP), which has been
funded by the National Astronomical Observatories of China, the Chinese
Academy of Sciences (the Strategic Priority Research Program “The Emergence of
Cosmological Structures” Grant XDB09000000), and the Special Fund for
Astronomy from the Ministry of Finance. The BASS is also supported by the
External Cooperation Program of Chinese Academy of Sciences (Grant
114A11KYSB20160057), and Chinese National Natural Science Foundation (Grant
12120101003, 11433005).

The Legacy Survey team makes use of data products from the Near-Earth Object
Wide-field Infrared Survey Explorer (NEOWISE), which is a project of the Jet
Propulsion Laboratory/California Institute of Technology. NEOWISE is funded by
the National Aeronautics and Space Administration.

The Legacy Surveys imaging of the DESI footprint is supported by the Director,
Office of Science, Office of High Energy Physics of the U.S. Department of
Energy under Contract No. DE-AC02-05CH1123, by the National Energy Research
Scientific Computing Center, a DOE Office of Science User Facility under the
same contract; and by the U.S. National Science Foundation, Division of
Astronomical Sciences under Contract No. AST-0950945 to NOAO.

This work has made use of data from the Asteroid Terrestrial-impact Last Alert
System (ATLAS) project. The Asteroid Terrestrial-impact Last Alert System
(ATLAS) project is primarily funded to search for near earth asteroids through
NASA grants NN12AR55G, 80NSSC18K0284, and 80NSSC18K1575; byproducts of the NEO
search include images and catalogs from the survey area. This work was
partially funded by Kepler/K2 grant J1944/80NSSC19K0112 and HST GO-15889, and
STFC grants ST/T000198/1 and ST/S006109/1. The ATLAS science products have
been made possible through the contributions of the University of Hawaii
Institute for Astronomy, the Queen’s University Belfast, the Space Telescope
Science Institute, the South African Astronomical Observatory, and The
Millennium Institute of Astrophysics (MAS), Chile.

This work was supported by the Australian government through the Australian
Research Council’s Discovery Projects funding scheme (DP200102471).

## Data Availability

A public release of the entire eRASS1 data taken within the German half of the
eROSITA sky is anticipated for Q4 2023. eRASS2-5 data are expected to be
released at a later stage. The Swift data is available to download through the
UK Swift Data Science website171717https://www.swift.ac.uk/archive/index.php.
The XMM data will become public after the propietory period expires
(2023-09-21). Publicly available ATLAS data can be accessed through the ATLAS
forced photometry service181818https://fallingstar-data.com/forcedphot/.
Publicly available ZTF data can be accessed through the ZTF forced photometry
service191919https://irsa.ipac.caltech.edu/Missions/ztf.html. NTT/EFOSC2
spectroscopy has been obtained under the program IDs 108.220C.012 (PI. C.
Inserra) and 109.23JL.001 (PI. I. Grotova). All optical spectra are publicly
available. ATCA data are stored in the Australia Telescope Online
Archive202020https://atoa.atnf.csiro.au/, and will become publicly accessible
18 months from the date of observation.

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##  Appendix A Additional photometric information

The optical and UV photometry of AT 2022dsb is presented in Table 9.

Table 9: Optical and UV photometry of AT 2022dsb, corrected for Galactic extinction, and with 3$\sigma$ upper limits presented. MJD | Instrument | Filter | Magnitude
---|---|---|---
59601.660 | ATLAS | o |  $<$19.21
59610.660 | ATLAS | o |  $<$18.97
59619.140 | ATLAS | c |  $<$18.14
59624.095 | ATLAS | o |  $<$18.69
59624.118 | ATLAS | o |  $<$19.59
59627.077 | ATLAS | o |  $<$18.90
59627.082 | ATLAS | o |  $<$18.90
59627.086 | ATLAS | o |  $<$19.01
59627.092 | ATLAS | o |  $<$18.97
59628.113 | ATLAS | o |  $<$18.84
59628.113 | ATLAS | o | 18.88 $\pm$ 0.34
59628.127 | ATLAS | o |  $<$19.13
59628.129 | ATLAS | o |  $<$18.97
59628.129 | ATLAS | o | 18.64 $\pm$ 0.26
59628.139 | ATLAS | o |  $<$18.58
59628.140 | ATLAS | o |  $<$18.46
59629.365 | ATLAS | o | 18.45 $\pm$ 0.13
59630.344 | ATLAS | o | 18.19 $\pm$ 0.13
59630.347 | ATLAS | o | 18.83 $\pm$ 0.20
59630.364 | ATLAS | o | 17.73 $\pm$ 0.09
59630.371 | ATLAS | o | 18.06 $\pm$ 0.12
59637.330 | ATLAS | o | 17.77 $\pm$ 0.04
59643.113 | Swift | UVW1 | 16.24 $\pm$ 0.05
59643.114 | Swift | U | 16.24 $\pm$ 0.06
59643.114 | Swift | B | 15.68 $\pm$ 0.06
59643.116 | Swift | UVW2 | 15.52 $\pm$ 0.05
59643.118 | Swift | V | 15.37 $\pm$ 0.07
59643.120 | Swift | UVM2 | 15.87 $\pm$ 0.05
59644.110 | ATLAS | o | 18.12 $\pm$ 0.05
59645.340 | ATLAS | c | 17.82 $\pm$ 0.06
59647.070 | ATLAS | o | 17.77 $\pm$ 0.05
59648.070 | ATLAS | o | 17.85 $\pm$ 0.03
59649.769 | Swift | UVW1 | 16.28 $\pm$ 0.06
59649.771 | Swift | U | 16.16 $\pm$ 0.06
59649.771 | Swift | B | 15.64 $\pm$ 0.06
59649.773 | Swift | UVW2 | 15.61 $\pm$ 0.05
59649.774 | Swift | V | 15.34 $\pm$ 0.07
59649.776 | Swift | UVM2 | 15.91 $\pm$ 0.05
59653.330 | ATLAS | o | 17.90 $\pm$ 0.04
59654.350 | ATLAS | o | 18.27 $\pm$ 0.06
59655.400 | ATLAS | o | 18.20 $\pm$ 0.08
59656.020 | ATLAS | o | 18.26 $\pm$ 0.09
59656.152 | Swift | UVW1 | 16.94 $\pm$ 0.08
59656.153 | Swift | U | 16.60 $\pm$ 0.09
59656.154 | Swift | B | 16.04 $\pm$ 0.08
59656.156 | Swift | UVW2 | 16.54 $\pm$ 0.07
59656.158 | Swift | V | 15.29 $\pm$ 0.08
59656.160 | Swift | UVM2 | 16.59 $\pm$ 0.08
59657.330 | ATLAS | o | 18.29 $\pm$ 0.08
59658.300 | ATLAS | o | 18.23 $\pm$ 0.09
59662.600 | ATLAS | o | 18.74 $\pm$ 0.25
59663.970 | Swift | UVW1 | 17.26 $\pm$ 0.12
59666.160 | ATLAS | o | 19.01 $\pm$ 0.12
59668.126 | Swift | UVW1 | 17.42 $\pm$ 0.09
59668.127 | Swift | U | 16.98 $\pm$ 0.09
59668.127 | Swift | B | 16.12 $\pm$ 0.07
59668.128 | Swift | UVW2 | 16.85 $\pm$ 0.07
59668.129 | Swift | V | 15.58 $\pm$ 0.08
59668.130 | Swift | UVM2 | 17.09 $\pm$ 0.10
59669.190 | ATLAS | o | 19.50 $\pm$ 0.15
59670.310 | ATLAS | o | 19.94 $\pm$ 0.23
59670.461 | ZTF | r | 19.10 $\pm$ 0.11
59670.484 | ZTF | g | 19.07 $\pm$ 0.07
59673.030 | ATLAS | c | 18.51 $\pm$ 0.10
59673.417 | ZTF | g | 19.08 $\pm$ 0.11
59673.481 | ZTF | r | 19.18 $\pm$ 0.17
59673.807 | Swift | UVW1 | 17.58 $\pm$ 0.11
59673.809 | Swift | U | 16.99 $\pm$ 0.10
59673.809 | Swift | B | 16.11 $\pm$ 0.08
59673.811 | Swift | UVW2 | 17.17 $\pm$ 0.12
59674.260 | ATLAS | o | 19.33 $\pm$ 0.12
59675.461 | ZTF | g | 19.15 $\pm$ 0.10
59677.314 | Swift | UVW1 | 17.84 $\pm$ 0.11
59677.315 | Swift | U | 17.19 $\pm$ 0.10
59677.316 | Swift | B | 16.15 $\pm$ 0.07
59677.317 | Swift | UVW2 | 17.26 $\pm$ 0.10
59677.360 | ATLAS | o | 19.04 $\pm$ 0.22
59677.462 | ZTF | g | 19.18 $\pm$ 0.14
59678.280 | ATLAS | o | 19.13 $\pm$ 0.12
59680.020 | ATLAS | o | 19.94 $\pm$ 0.17
59682.270 | ATLAS | o |  $<$20.34
59683.040 | ATLAS | o |  $<$17.64
59683.207 | Swift | UVW1 | 17.81 $\pm$ 0.16
59683.208 | Swift | U | 17.34 $\pm$ 0.16
59683.208 | Swift | B | 16.15 $\pm$ 0.10
59683.210 | Swift | UVW2 | 17.58 $\pm$ 0.15
59683.211 | Swift | V | 15.59 $\pm$ 0.12
59683.213 | Swift | UVM2 | 17.47 $\pm$ 0.13
59683.440 | ZTF | g | 19.55 $\pm$ 0.26
59684.416 | ZTF | r |  $<$19.43
59684.451 | ZTF | g | 19.38 $\pm$ 0.28
59685.190 | ATLAS | o | 19.65 $\pm$ 0.34
59689.418 | ZTF | r |  $<$19.06
59689.461 | ZTF | g |  $<$19.06
59690.320 | ATLAS | o | 19.38 $\pm$ 0.24
59693.280 | ATLAS | o |  $<$20.41
59694.330 | ATLAS | o |  $<$20.21
59694.377 | ZTF | g |  $<$20.38
59696.337 | ZTF | g | 19.92 $\pm$ 0.21
59696.447 | ZTF | r | 20.04 $\pm$ 0.30
59696.480 | Swift | UVM2 | 17.74 $\pm$ 0.16
59696.482 | Swift | UVW1 | 18.24 $\pm$ 0.17
59696.482 | Swift | U | 17.34 $\pm$ 0.11
59696.483 | Swift | UVW2 | 18.15 $\pm$ 0.18
59697.300 | ATLAS | o | 19.55 $\pm$ 0.14
59698.270 | ATLAS | o | 19.62 $\pm$ 0.17
59699.020 | ATLAS | o | 20.06 $\pm$ 0.22
59699.355 | ZTF | r |  $<$20.37
59699.440 | ZTF | g | 19.80 $\pm$ 0.17
59700.070 | ATLAS | o |  $<$20.69
59701.270 | ATLAS | o |  $<$18.80
59701.376 | ZTF | g | 20.46 $\pm$ 0.31
59701.418 | ZTF | g | 19.89 $\pm$ 0.17
59701.424 | Swift | UVM2 | 18.61 $\pm$ 0.27
59701.425 | Swift | UVW1 | 18.17 $\pm$ 0.16
59701.426 | Swift | U | 17.41 $\pm$ 0.11
59701.427 | Swift | UVW2 | 17.76 $\pm$ 0.13
59702.290 | ATLAS | o |  $<$20.37
59703.470 | ATLAS | o |  $<$20.87
59704.748 | Swift | UVM2 | 17.90 $\pm$ 0.14
59704.750 | Swift | UVW1 | 18.16 $\pm$ 0.13
59704.752 | Swift | U | 17.48 $\pm$ 0.09
59704.753 | Swift | UVW2 | 18.19 $\pm$ 0.15
59705.310 | ATLAS | o |  $<$18.27
59706.240 | ATLAS | o |  $<$20.64
59707.397 | ZTF | r |  $<$20.14
59709.230 | ATLAS | o | 19.69 $\pm$ 0.21
59709.378 | ZTF | r |  $<$20.13
59710.310 | ATLAS | o |  $<$17.74
59712.540 | ATLAS | o |  $<$18.81
59723.990 | ATLAS | o | 19.78 $\pm$ 0.25
59725.200 | ATLAS | o | 19.66 $\pm$ 0.12
59726.590 | ATLAS | o | 19.85 $\pm$ 0.15
59729.054 | Swift | U | 17.24 $\pm$ 0.16
59729.055 | Swift | UVW2 | 18.19 $\pm$ 0.24
59729.577 | Swift | UVW1 | 18.15 $\pm$ 0.21
59729.978 | Swift | UVM2 | 18.38 $\pm$ 0.20
59730.920 | ATLAS | o |  $<$20.52
59731.960 | ATLAS | o | 20.09 $\pm$ 0.24
59734.960 | ATLAS | o | 20.40 $\pm$ 0.36
59735.804 | Swift | UVM2 | 18.46 $\pm$ 0.22
59735.808 | Swift | UVW1 | 18.52 $\pm$ 0.17
59735.810 | Swift | U | 17.47 $\pm$ 0.10
59735.814 | Swift | UVW2 | 18.37 $\pm$ 0.15
59735.910 | ATLAS | o | 19.72 $\pm$ 0.14
59737.180 | ATLAS | o |  $<$20.72
59745.310 | ATLAS | o |  $<$18.20
59749.140 | ATLAS | o |  $<$20.22
59750.190 | ATLAS | o |  $<$19.89
59765.100 | ATLAS | o |  $<$19.91
59766.790 | ATLAS | o |  $<$20.59
59767.820 | ATLAS | o |  $<$20.07
59773.150 | ATLAS | o | 19.15 $\pm$ 0.34
59775.840 | ATLAS | o |  $<$20.29
59781.070 | ATLAS | o | 19.45 $\pm$ 0.18
59782.100 | ATLAS | o |  $<$20.57
59783.840 | ATLAS | c |  $<$20.43
59785.100 | ATLAS | o |  $<$20.55
59786.060 | ATLAS | o |  $<$20.03
59787.860 | ATLAS | c | 20.58 $\pm$ 0.36
59789.060 | ATLAS | c | 20.40 $\pm$ 0.28
59790.410 | ATLAS | c | 20.53 $\pm$ 0.29
59791.780 | ATLAS | c |  $<$20.71
59793.080 | ATLAS | c | 20.31 $\pm$ 0.27
59795.830 | ATLAS | o |  $<$19.88
59799.810 | ATLAS | o |  $<$16.37
59801.080 | ATLAS | o |  $<$19.45
59802.070 | ATLAS | o |  $<$17.68
59805.020 | ATLAS | o |  $<$20.24
59806.450 | ATLAS | o |  $<$20.44
59807.750 | ATLAS | o |  $<$20.52
59814.790 | ATLAS | o |  $<$20.43
59815.740 | ATLAS | o | 19.70 $\pm$ 0.35
59817.520 | ATLAS | c | 19.96 $\pm$ 0.15
59818.010 | ATLAS | c |  $<$19.87
59818.750 | ATLAS | o |  $<$20.19
59820.990 | ATLAS | c | 19.22 $\pm$ 0.17
59821.010 | ATLAS | c | 20.30 $\pm$ 0.32
59822.050 | ATLAS | c | 20.03 $\pm$ 0.25
59822.790 | ATLAS | o |  $<$18.97
59825.980 | ATLAS | o |  $<$19.02
59828.990 | ATLAS | o |  $<$19.25
59830.770 | ATLAS | o |  $<$19.78
59855.514 | Swift | UVW2 | 18.53 $\pm$ 0.15
59855.518 | Swift | UVM2 | 18.52 $\pm$ 0.17
59855.524 | Swift | UVW1 | 18.61 $\pm$ 0.13

##  Appendix B Additional spectroscopic information

A zoom-in on the evolution of the broad He ii and H$\alpha$ line profiles is
presented in Fig. 15, whilst details on the spectroscopic observations and
data reduction are presented below (see Table 10 for an observation log). The
archival 6dFGS optical spectrum taken in 2002 is plotted in Fig. 16.

Figure 15: Evolution of the H$\alpha$ (top) and He ii (bottom) line regions for the first four spectra that clearly show broad emission lines. The vertical lines mark [N ii] 6548Å, H$\alpha$, and [N ii] 6583Å in the top panel (left to right), and N iii 4640Å, He ii 4686Å, H$\beta$, [O iii] 4959Å and [O iii] 5007Å in the bottom panel. Table 10: Spectroscopic observations of AT 2022dsb. MJD | UT | Telescope | Instrument | Airmass | Exposure [s]
---|---|---|---|---|---
59636 | 2022-02-26 | LCO FTN | FLOYDS | 1.4 | 2400
59640 | 2022-03-02 | NTT | EFOSC2 | 1.0 | 600
59665 | 2022-03-27 | SALT | RSS | 1.3 | 500
59680 | 2022-04-11 | SALT | RSS | 1.2 | 500
59680 | 2022-04-11 | LCO FTS | FLOYDS | 1.1 | 2400
59699 | 2022-04-30 | LCO FTS | FLOYDS | 1.2 | 2400
59726 | 2022-05-27 | LCO FTN | FLOYDS | 1.5 | 2400
59767 | 2022-07-07 | NTT | EFOSC2 | 1.2 | 1800
59781 | 2022-07-21 | LCO FTS | FLOYDS | 1.4 | 2400
Figure 16: Archival optical spectrum of the host galaxy of AT 2022dsb taken in
2002 by 6dFGS (Jones et al., 2009). The broad He ii emission complex around
4640 Åseen in the optical specroscopic follow-up campaign of AT 2022dsb is
clearly not detected. The publicly available 6dFGS spectra do not have
absolute flux calibrations.

FLOYDS: AT 2022dsb was observed five times with the FLOYDS spectrographs
mounted at the Las Cumbres Observatory 2 m telescopes at Siding Springs
Observatory (FTS, Australia) and Haleakala (FTN, Hawaii). Observations were
performed on 2022-2-26 (FTN), 2022-4-11 and 2022-4-30 (both FTS), 2022-5-27
(FTN), and 2022-7-21 (FTS) with 2400 s exposure each. The two identical FLOYDS
are cross-dispersed, low-resolution spectrographs covering the wavelength
range from 320 to 1000 nm with a resolution between R=400 and R=700. Data were
automatically processed by the standard FLOYDS pipeline.

SALT: Spectroscopy of AT 2022dsb was undertaken with the Southern African
Large Telescope (SALT; Buckley et al. 2006) using the Robert Stobie
Spectrograph (RSS, Burgh et al. 2003), on two nights: 2022 March 27 and 2022
April 11, starting 23:02 and 03:28 UTC, respectively. Two consecutive 500 sec
exposure spectra with the PG0900 grating were obtained on each night. The
spectra covered the region 4345$-$7400Åat a mean resolution of 5.7Å. The
spectra were reduced using the PyRAF-based PySALT package (Crawford et al.
2010)212121https://astronomers.salt.ac.za/software/, which includes
corrections for gain and cross-talk, and performs bias subtraction. We
extracted the science spectrum using standard IRAF222222https://iraf-
community.github.io/ tasks, including wavelength calibration (Xenon
calibration lamp exposures were taken, immediately after the science spectra),
background subtraction, and 1D spectra extraction. Due to SALT’s optical
design, absolute flux calibration is not possible232323The entrance pupil
changes its position during an observation, resulting in a changing effective
collecting area.. Observations of spectrophotometric standards during twilight
were used to obtain relative flux calibration.

EFOSC2: The source was re-observed with the EFOSC2 spectrograph mounted at the
ESO New Technology Telescope (NTT) at La Silla observatory on 2022-7-7 as part
of a program targeting eROSITA-selected TDEs (PI: Grotova; 109.23JL.001).
Observations were performed with Grism#13 covering approx. 3600-9300 Å with a
resolution of $\approx 21$ Å, and an exposure time of 1800s. The data were
standard reduced with the IRAF (Tody, 1986) community
distribution242424https://iraf-community.github.io. The EFOSC2 spectrum taken
on 2022-03-02 also used Grism#13, and was reduced by the ePESSTO+ team (Smartt
et al., 2015) using version ntt_2.4.0 of their NTT reduction
pipeline252525https://github.com/svalenti/pessto; the reduced spectrum was
downloaded from the Weizmann Interactive Supernova Data Repository
(WISeREP262626https://www.wiserep.org/).