# To TDE or not to TDE: The luminous transient ASASSN-18jd with TDE-like and
AGN-like qualities

J. M. M. Neustadt,1 T. W.-S. Holoien,2 C. S. Kochanek,1,3 K. Auchettl,4,5 J.
S. Brown,5 B. J. Shappee,6 R. W. Pogge,1 Subo Dong,7 K. Z. Stanek,1 M. A.
Tucker,6 S. Bose,1,3 Ping Chen,7 C. Ricci,8,7 P. J. Vallely,1 J. L. Prieto,8,9
D. A. Coulter,5 M. R. Drout,10 R. J. Foley,5 C. D. Kilpatrick,5 A. L. Piro,2
C. Rojas-Bravo,5 D. A. H. Buckley,11 M. Gromadzki,12 G. Dimitriadis,5 M. R.
Siebert,5 A. Do,6 M. E. Huber,6 A. V. Payne6
1Department of Astronomy, The Ohio State University, 140 West 18th Avenue,
Columbus, OH 43210, USA
2The Observatories of the Carnegie Institution for Science, 813 Santa Barbara
Street, Pasadena, CA 91101, USA
3Center for Cosmology and AstroParticle Physics (CCAPP), The Ohio State
University, 191 W. Woodruff Avenue, Columbus, OH 43210, USA
4DARK, Niels Bohr Institute, University of Copenhagen, Lyngbyvej 2, 2100
Copenhagen, Denmark
5Department of Astronomy and Astrophysics, University of California, Santa
Cruz, CA 95064, USA
6Institute for Astronomy, University of Hawai’i, 2680 Woodlawn Drive,
Honolulu, HI 96822, USA
7Kavli Institute for Astronomy and Astrophysics, Peking University, Yi He Huan
Road 5, Hai Dian District, Beijing 100871, China
8Núcleo de Astronomía de la Facultad de Ingeniería y Ciencias, Universidad
Diego Portales, Av. Ejército 441 Santiago, Chile
9Millennium Institute of Astrophysics, Santiago, Chile
10Department of Astronomy and Astrophysics, University of Toronto, 50 St.
George St., Toronto, ON M5S 3H4, Canada
11South African Astronomical Observatory, PO Box 9, Observatory 7935, Cape
Town, South Africa
12Astronomical Observatory, University of Warsaw, Al. Ujazdowskie 4, 00-478
Warszawa, Poland E-mail<EMAIL_ADDRESS>(JMMN)

(Accepted XXX. Received YYY; in original form ZZZ)

###### Abstract

We present the discovery of ASASSN-18jd (AT 2018bcb), a luminous
optical/UV/X-ray transient located in the nucleus of the galaxy 2MASX
J22434289–1659083 at $z=0.1192$. Swift UVOT photometry shows the UV SED of the
transient to be well modeled by a slowly shrinking blackbody with temperature
$T\sim 2.5\times 10^{4}\rm~{}K$, a maximum observed luminosity of
$L_{\text{max}}=4.5^{+0.6}_{-0.3}\times 10^{44}\rm~{}erg~{}s^{-1}$, and a
total radiated energy of $E=9.6^{+1.1}_{-0.6}\times 10^{51}\rm~{}erg$. X-ray
data from Swift XRT and XMM-Newton show a transient, variable X-ray flux with
blackbody and power-law components. Optical spectra show strong, roughly
constant broad Balmer emission as well as transient features attributable to
He ii, N iii-v, O iii, and coronal Fe. While ASASSN-18jd shares similarities
with Tidal Disruption Events (TDEs), it is also similar to the “rapid turn-on”
events seen in quiescent galaxies and in faint Active Galactic Nuclei (AGNs).

###### keywords:

accretion, accretion discs – black hole physics – galaxies: nuclei.

††pubyear: 2019††pagerange: To TDE or not to TDE: The luminous transient
ASASSN-18jd with TDE-like and AGN-like qualities–6

##  1 Introduction

Active galactic nuclei (AGNs) are known to vary, both photometrically and
spectroscopically. These variations are likely driven by the variable
accretion of material onto the Supermassive Black Hole (SMBH) at the center of
the galaxy, but there are contributions from variable obscuration of
components of the AGN by gas or dust. While these variations have been studied
for decades (e.g., Andrillat 1968; Tohline & Osterbrock 1976; Oknyanskij
1978), recent studies have continued to show a diversity of events that are
qualitatively different from the typical, modest amplitude, stochastic
variability observed in all AGNs (e.g., MacLeod et al. 2012). In particular,
there are changing-look AGNs (e.g., Bianchi et al. 2005; Shappee et al. 2014;
MacLeod et al. 2016), where broad emission lines appear and/or disappear over
timescales of years, and “rapid turn-on” events (e.g. Gezari et al. 2017a;
Tadhunter et al. 2017; Gromadzki et al. 2019; Trakhtenbrot et al. 2019;
Frederick et al. 2019), where a blue continuum and broad emission lines appear
over timescales of months.

There are also rapid, luminous flares occurring in galaxies with previously
quiescent galactic nuclei or faint AGNs that are believed to be Tidal
Disruption Events (TDEs). TDEs are the result of a star crossing the tidal
radius of the SMBH and being disrupted by tidal forces from the SMBH (Rees,
1988; Phinney, 1989; Evans & Kochanek, 1989). Some of the disrupted material
is then accreted onto the SMBH, producing a bright optical/UV/X-ray transient
that fades over time. The relative amounts of accreted and ejected material,
as well as the properties of the transient, are a combination of various
physical processes, such as geometry of the original orbit (e.g., Guillochon &
Ramirez-Ruiz 2013; Dai et al. 2018), the properties of both the disrupted star
and SMBH (e.g., Guillochon & Ramirez-Ruiz 2013; Kochanek 2016), and radiative
feedback from the accretion (e.g., Gaskell & Rojas Lobos 2014; Strubbe &
Murray 2015; Roth et al. 2016; Roth & Kasen 2018). Previous studies of TDEs
(e.g., Holoien et al. 2014; Auchettl et al. 2017) generally found that the
energy radiated during the events are of order $10^{51}\rm~{}erg$ with mass
equivalents of less than $0.01\rm~{}M_{\odot}$ assuming an accretion
efficiency of $\eta=0.1$, suggesting that most of the bound debris is ejected
and not accreted, or that the accretion efficiency is lower than expected.
However, when one considers the energy emitted over much longer timescales,
then the total radiated energy approaches $\sim$0.1 M⊙ (van Velzen et al.,
2019b).

Compared to other extragalactic transients, such as supernovae (SNe) and AGN
variability, TDEs have unique properties (see, e.g., Hung et al. 2017;
Auchettl et al. 2018; Holoien et al. 2019a). The UV/optical spectral energy
distributions (SEDs) of TDEs are well modeled by blockbodies with temperatures
of a few $10^{4}$ K. This is initially true for SNe, but SNe rapidly cool to
below $10^{4}$ K within weeks, whereas TDEs do not (see Figure 10 from Holoien
et al. 2019a). The UV/optical SEDs of AGNs are usually best fit with power
laws, $f_{\lambda}\propto\lambda^{-\alpha}$, rather than a single-temperature
blackbody. A single-epoch AGN spectrum usually has $1<\alpha<2$ (Koratkar &
Blaes, 1999; Vanden Berk et al., 2001), whereas an AGN difference spectrum
(i.e. the difference between bright and faint epochs) often has $\alpha\gtrsim
2$ (see, e.g., Ruan et al. 2014; Hung et al. 2016), thus becoming “bluer when
brighter” (Wilhite et al., 2005). This is in line with theoretical predictions
from a thin accretion disc model with a $T\propto R^{3/4}$ temperature
profile, producing an SED with $\alpha\simeq 2.3$ in the optical/near UV
(Shakura & Sunyaev, 1973).

The light curves of TDEs usually peak and decay monotonically, but several
TDEs show deviations from a monotonic decay. Examples are PS18kh/AT 2018zr
(Holoien et al., 2019a; van Velzen et al., 2019a), which rebrightened multiple
times after its initial decay, ASASSN-18ul/AT 2018fyk (Wevers et al., 2019b),
which showed an extended plateau in its light curve, and ASASSN-19bt (Holoien
et al., 2019b), which showed a short flare $\sim$30 d before reaching peak
brightness. Additionally, there is some evidence that TDEs may fade more
slowly as the SMBH mass increases (Blagorodnova et al., 2017; Wevers et al.,
2017; van Velzen et al., 2019b). By contrast, AGN variability consists of
stochastic, tenths-of-a-magnitude variation over periods of tens to hundreds
of days (MacLeod et al., 2012).

Another distinction between TDEs and AGNs is their X-ray emission. TDEs
usually have smaller column densities ($N_{\text{H}}$) and intrinsically
softer spectra (e.g., ASASSN-14li: Brown et al. 2017; Kara et al. 2018) than
that of AGNs. While TDEs show little variation in their hardness ratios as
they fade, AGNs tend to become softer as they brighten and become harder as
they fade (e.g., Auchettl et al. 2018).

The optical spectra of TDEs are usually dominated by very broad (FWHM $\gtrsim
10^{4}\rm~{}km~{}s^{-1}$) H and/or He ii lines (e.g., Arcavi et al. 2014).
Usually, the relative strengths of H to He ii emission in an individual TDE
are relatively constant, although there are counter examples: ASASSN-15oi had
weak, transient H features (Holoien et al., 2018), while AT 2017eqx
transitioned from being H-dominated to being He ii-dominated (Nicholl et al.,
2019). Many TDEs also show optical emission lines attributable to N iii and O
iii (Blagorodnova et al., 2019; Leloudas et al., 2019). By contrast, optical
spectra of Type I AGNs are dominated by broad H emission (FWHM $\sim
2000\rm~{}km~{}s^{-1}$) and narrow, collisionally excited lines like [O iii]
and [N ii]. For Type II AGNs, the H emission lines are also narrow. The broad
and narrow lines are thought to originate from different locations in the AGN,
with the Broad Line Region (BLR) being closer to the SMBH than the Narrow Line
Region (NLR) (Peterson, 1993). AGNs will sometimes show weak He i emission and
even weaker He ii emission (Vanden Berk et al., 2001). While TDE emission
lines generally become narrower as the continuum fades (e.g., Holoien et al.
2016a), the broad lines of AGNs show the opposite behavior, as predicted by
photoionization models (Peterson et al., 2004; Denney et al., 2009).

Here we discuss ASASSN-18jd, a transient source first detected by the All-Sky
Automated Survey for Supernovae (ASAS-SN, Shappee et al. 2014; Kochanek 2016)
on 2018-04-09 (MJD 58217.4) at ($\alpha,\delta$) $=$ (22:43:42.866,
–16:59:08.410). The event was coincident with the center of the galaxy 2MASX
J22434289–1659083. We first reported the discovery of the transient to the
Transient Name Server (TNS) on 2018-04-27 (MJD 58235.1) (Bersier & Stanek,
2018), where it was designated AT 2018bcb. ASASSN-18jd was discovered after a
seasonal gap due to Sun-constraints, and the last observation before the
seasonal gap was on 2017-12-14 (MJD 58101.0), $-$116.4 d before detection.
ASASSN-18jd was fading when discovered, so the peak brightness probably
occurred during the gap. Throughout this paper, we use the date of first
detection in the ASAS-SN data, MJD 58217.4, as a reference date for tracking
the evolution of the transient. Using the H$\alpha$ emission features in our
spectra (see Section 5), we find a redshift $z=0.1192$ for ASASSN-18jd. This
corresponds to a luminosity distance $D_{\text{L}}=559.6$ Mpc for a flat
universe with $h=0.696$, $\Omega_{\text{M}}=0.286$, and
$\Omega_{\Lambda}=0.714$ (Wright, 2006). The Galactic extinction along this
line-of-sight is $A_{V}=0.098$ mag (Schlafly & Finkbeiner, 2011). When we
correct for extinction, we use a Cardelli et al. (1989) extinction curve with
$R_{V}=3.1$ throughout.

In Section 2, we describe the archival pre-outburst photometry of the host
galaxy as well as the new photometry and spectroscopy for ASASSN-18jd. In
Sections 3, 4, 5, and 6, we use this data to characterize the UV/optical SED,
the X-ray properties, the optical spectra, and the UV spectra of ASASSN-18jd,
respectively. We compare these properties to well-studied TDEs and to the
typical properties of AGNs. In Section 7, we place ASASSN-18jd within the
context of TDEs as well as other SMBH-driven transients like changing-look
AGNs and the newly-identified rapid turn-on events.

##  2 Observations

In this section, we summarize the archival data available for the host galaxy
and our new photometry and spectroscopy of ASASSN-18jd.

###  2.1 Host photometry

The available UV, optical, and IR photometry of the host galaxy 2MASX
J22434289–1659083 is summarized in Table 1. The data come from the Galaxy
Evolution Explorer (GALEX, Martin et al. 2005; Bianchi et al. 2011), Data
Release 1 of Pan-STARRS (Pan-STARRS1, Chambers et al. 2016; Flewelling et al.
2016), the Two Micron All-Sky Survey (2MASS, Skrutskie et al. 2006) Extended
Survey Catalog, and the Wide Infrared Survey Explorer (WISE, Wright et al.
2010). Due to the compactness of the galaxy, we used a region of
$5^{\prime\prime}$ to derive the SED of the host, as this captures most of the
emission of the galaxy in the relevant wavelengths. There were no archival
observations available from the Hubble Space Telescope (HST), the Chandra
X-ray Observatory, the X-Ray Multi-Mirror Mission (XMM-Newton), or the Dark
Energy Survey (DES) Data Release 1. Due to its location in the southern sky,
the host was not observed as part of the Sloan Digital Sky Survey (SDSS) or
the VLA FIRST Survey. There are no previous outbursts or signs of $V$-band
variability from the host galaxy between December 2005 and July 2010 in the 11
epochs of data from the Catalina Real-time Transient Survey (CRTS, Drake et
al. 2009). The host has a WISE color of $(\textit{W}1-\textit{W}2)=0.28\pm
0.04$ mag (Vega) which is bluer than most luminous AGNs with
$(\textit{W}1-\textit{W}2)\geq 0.8$ mag (Stern et al., 2012; Assef et al.,
2013).

Table 1: Observed (left) and FAST-derived (right) $5^{\prime\prime}$ aperture AB magnitudes of host galaxy 2MASX J22434289–1659083. Errors are assumed to be 0.1 mag (0.2 mag for GALEX). Filter | Observed | Synthetic
---|---|---
GALEX NUV | 21.11 | 20.98
PS $g$ | 17.77 | 17.83
PS $r$ | 16.99 | 16.97
PS $i$ | 16.55 | 16.57
PS $z$ | 16.34 | 16.28
PS $y$ | 16.10 | 16.10
WISE $W1$ | 16.16 | $-$
WISE $W2$ | 16.53 | $-$
Swift UVW2 | $-$ | 20.99
Swift UVM2 | $-$ | 21.01
Swift UVW1 | $-$ | 20.82
Swift U | $-$ | 19.77
Swift B | $-$ | 18.54
Swift V | $-$ | 17.33
SDSS $u^{\prime}$ | $-$ | 19.61
SDSS $g^{\prime}$ | $-$ | 17.98
SDSS $r^{\prime}$ | $-$ | 16.96
SDSS $i^{\prime}$ | $-$ | 16.54

There is no source coincident with the host galaxy in the ROSAT All-Sky Survey
(Voges et al., 1999; Boller et al., 2016). Assuming a power law of
$\Gamma=1.75$ and using the line-of-sight Galactic H i column density of
$N_{\text{H}}=2.71\times 10^{20}\rm~{}cm^{-2}$ (Kalberla et al., 2005), this
yields a $3\sigma$-upper limit on the X-ray flux in the 0.3$-$10.0 keV band of
$F_{0.3-10}\lesssim 2\times 10^{-12}\rm~{}erg~{}s^{-1}~{}cm^{-2}$,
corresponding to a luminosity of $L_{0.3-10}\lesssim 7.5\times
10^{43}\rm~{}erg~{}s^{-1}$ and $L_{X}/L_{\text{Edd}}<0.015$ given our estimate
of the SMBH mass (computed below). This limit does not exclude weak AGN
activity, as surveys have found Type I and II AGNs at and below this X-ray
luminosity (Tozzi et al., 2006; Marchesi et al., 2016; Liu et al., 2016; Ricci
et al., 2017). There is also no source coincident with the host galaxy in the
1.4 GHz NRAO VLA Sky Survey (NVSS, Condon et al. 1998). This limits the flux
density at 1.4 GHz to be $S(1.4\rm~{}GHz)\leq 2.5~{}mJy$, corresponding to
$L_{\nu}\rm(1.4~{}GHz)\leq 9.4\times 10^{29}~{}erg~{}s^{-1}~{}Hz^{-1}$. All of
these factors imply that any pre-event AGN activity must be relatively weak,
if present at all.

We used the archival GALEX, Pan-STARRS1, and WISE fluxes to model the host
galaxy’s SED with the code Fitting and Assessment of Synthetic Templates
(FAST; Kriek et al. 2009). In doing the fits, we assumed minimum errors of 0.1
mag to account for potential systematic errors, except for the GALEX data,
where we used the reported errors of 0.2 mag. We included Galactic extinction,
an exponentially declining star-formation history, a Salpeter initial mass
function, and the Bruzual & Charlot (2003) stellar population models. Our best
fit had a stellar mass of $M_{*}=1.7^{+0.1}_{-0.9}\times
10^{11}\rm~{}M_{\odot}$, age = $8.9^{+0.6}_{-5.9}$ Gyr, and a star formation
rate of SFR $=0.6^{+0.1}_{-0.3}\rm~{}M_{\odot}~{}yr^{-1}$. We combined the SED
fit generated by FAST and the appropriate filter profiles to produce our
synthetic photometry. Table 1 compares these FAST-derived magnitudes to the
data and provides estimates of the host flux in the Swift bands.

We used the stacked Pan-STARRS1 $g$\- and $r$-band images and the galaxy image
decomposition program GALFIT (Peng et al., 2002) to estimate the bulge-to-
total light ratio, $B/T$. We found $(B/T)_{g}\approx 0.18\pm 0.05$ and
$(B/T)_{r}\approx 0.16\pm 0.05$. Using these ratios, we assumed
$(B/T)_{V}\approx 0.17\pm 0.10$, corresponding to a bulge luminosity of
$\log{L_{V}/\rm L_{\odot}}=9.6\pm 0.3$ using the synthetic $V$-band luminosity
from FAST. This was done after correcting for the
$5\aas@@fstack{\prime\prime}{0}$ aperture of the synthetic FAST flux, but
because the galaxy is relatively compact, this correction changed $B/T$ by
only $\sim$0.01, well within the uncertainties. Using the
$M_{\text{BH}}$$-$$L_{V}$ relation from McConnell & Ma (2013), the SMBH mass
is then $\log{M_{\text{BH}}/\rm M_{\odot}}=7.6\pm 0.4$. The dispersion in the
$M_{\text{BH}}$$-$$L_{V}$ relation is much larger than any reasonable
uncertainties in $B/T$. While this mass is significantly larger than most SMBH
masses associated with TDEs (Wevers et al., 2017, 2019a), it is still in the
mass range where main sequence stars can be disrupted before crossing the
event horizon (see, e.g., Kochanek 2016).

###  2.2 ASAS-SN photometry

Figure 1: ASAS-SN photometry of ASASSN-18jd in $g$\- (blue) and $V$\- (red)
band magnitudes. Detections are presented as circles, and 3$\sigma$ upper
limits are triangles. The transient photometry excludes the flux of the host,
and the pre-transient “detections” are almost certainly artefacts of image
subtraction. The $3\sigma$ upper limit on the average pre-transient $g$-band
variability for Nov and Dec 2017 (MJD $\sim$ 58060$-$58100) is also shown. The
seasonal gaps, where no data were collected in epochs relevant to the
transient, are shaded.

We show the full ASAS-SN $g$\- and $V$-band light curve before and during the
transient in Figure 1, where the flux of the host galaxy has been subtracted.
We also show the upper limit on the pre-transient $g$-band flux from November
and December 2017 of $g>20.09\rm~{}mag$. The $V$-band photometry from November
2013 to December 2017 shows no obvious variability, similar to what we see in
the CRTS data. Some epochs show faint “detections” of transient flux from the
host, but these are very likely artefacts from image subtraction, as none of
these “detections” are consistent across nearby epochs. The last pre-transient
epochs of observation for $g$ and $V$ band were 2017-11-29 (MJD 58086.1) and
2017-12-14 (MJD 58101.0), respectively. We use the latter date, MJD 58101.0,
116.4 d before the date of first detection, as the beginning of the seasonal
gap, which we later use to constrain the evolution of the transient in Section
3. Because ASAS-SN stopped using $V$ band in late 2018, there is no $V$-band
photometry after MJD 58447.2.

###  2.3 Swift photometry

We monitored ASASSN-18jd (Target ID: 10680; PI: T. Holoien) with the Swift
UVOT telescope (Roming et al., 2005) in the V (5468 Å), B (4392 Å), U (3465
Å), UVW1 (2600 Å), UVM2 (2246 Å), and UVW2 (1928 Å) filters. Each epoch of
UVOT data consists of 2 observations in each filter which we combined. We then
extracted counts from a $5\aas@@fstack{\prime\prime}{0}$ radius around the
source and use the background counts from a region with no sources and a
radius $\sim$40′′. For these two steps we use the HEAsoft software tasks
uvotimsum and uvotsource, respectively. The counts were then converted to AB
magnitudes and fluxes using the most recent UVOT calibrations (Poole et al.,
2008; Breeveld et al., 2010). Altogether, we obtained 50 epochs of Swift
observations from MJD 58249.8 (+32 d) to MJD 58661.4 (+444 d). There is a gap
in the Swift data from MJD 58489 (early January 2019) to MJD 58598 (mid April
2019) due to ASASSN-18jd becoming Sun-constrained.

Simultaneous with our Swift UVOT observations, we observed ASASSN-18jd in
photon-counting (PC) mode with the Swift X-ray Telescope (XRT, Burrows et al.
2005). All observations were reduced following the standard _Swift_ XRT data
reduction guide111http://swift.gsfc.nasa.gov/analysis/xrt_swguide_v1_2.pdf,
and reprocessed using the _Swift_ xrtpipeline version 0.13.2 script. Standard
filters and screening were applied, along with the most up-to-date calibration
files. We used a source region centered on the position of ASASSN-18jd with a
radius of $30^{\prime\prime}$ and a source free background region centered at
($\alpha,\delta$) $=$ (22:43:50.482,–16:54:32.90) with a radius of
$200^{\prime\prime}$. We corrected our count rate for the encircled energy
fraction ($\sim$90 per cent at 1.5 keV for a $30^{\prime\prime}$ radius, Hill
et al. 2004). We were also able to combine our Swift XRT observations into two
time bins covering the first and last $\sim$200 d of our monitoring program to
extract both a source and background spectrum. We merged our observations
using xselect version 12.9.1c, and then use the task xrtproducts and the same
extraction regions. Ancillary response files for each spectra are generated
using the task xrtmkarf, along with the standard response matrix files from
CALDB. A summary of these observations is included in Table 4.

###  2.4 XMM-Newton observations

Because ASASSN-18jd was X-ray bright, we requested two deep XMM-Newton
Observatory target of opportunity observations of the source. The first
observation was taken on 2018-05-28 (MJD 58266.8; ObsID: 0830191201, PI:
Schartel/Ricci), approximately +49 d after the initial discovery, while our
second observation was taken two weeks later on 2018-06-11 (MJD 58280.6;
ObsID: 0830191301, PI: Schartel/Ricci), approximately +63 d after the initial
discovery. Both the MOS and PN detectors were used for this analysis and both
observations were obtained in full frame mode using a thin filter. All data
reduction and analysis was done using the XMM-Newton science system (SAS)
version 17.0.0 (Gabriel et al., 2004) with the most up to date calibration
files.

To check for periods of high background activity that may affect the quality
of the data, we generated count rate histograms of the events that have
energies between 10$-$12 keV for each observation. We found that both
observations are only slightly affected by background flares, giving effective
exposure times of 32 ks, 32 ks, and 28 ks for the MOS1, MOS2 and PN detectors,
respectively for the first observation and 24 ks, 24 ks, and 21 ks,
respectively, for the second observation. A summary of these observations is
included in Table 4.

For our analysis we used the standard screening of events as specified in the
XMM-Newton analysis guide for the MOS and PN detectors222https://xmm-
tools.cosmos.esa.int/external/xmm_user_support/documentation/sas_usg/USG/. We
also corrected all event files for possible vignetting using the task
evigweight333https://xmm-
tools.cosmos.esa.int/external/sas/current/doc/evigweight. We extracted spectra
from both the MOS and PN detectors of ASASSN-18jd using the SAS task evselect
and the cleaned event files from all detectors. We used the same source region
that was used to analyse the Swift observations with a radius of
$30^{\prime\prime}$, while we used a source free background region centered at
($\alpha,\delta$) $=$ (20:43:25.495,–16:58:58.60) with a radius of
$70^{\prime\prime}$. All spectra were binned with a minimum of 20 counts per
energy bin using the ftools command grppha, and we used the X-ray spectral
fitting package (XSPEC) version 12.10.0 (Arnaud, 1996) and chi-squared
statistics to analyse the spectra. Count rates were also extracted from the PN
observations using the same regions and corrected for encircled energy
fraction444https://heasarc.nasa.gov/docs/xmm/uhb/onaxisxraypsf.html. These
data and their analysis are further discussed in Section 4.

###  2.5 Ground-based optical photometry

We obtained photometric observations from multiple ground-based observatories.
$BVgri$ observations were obtained using the Las Cumbres Observatory (Brown et
al., 2013) 1-m telescopes located at the Cerro Tololo Inter-American
Observatory (CTIO) in Chile, McDonald Observatory in Texas, Siding Spring
Observatory in Australia, and the South African Astronomical Observatory
(SAAO). $uBVgri$ observations were obtained with the Swope 1-m telescope at
Las Campanas Observatory in Chile. $BV$ observations were obtained with A
Novel Dual Imaging CAMera (ANDICAM; DePoy et al. 2003) on the SMARTS 1.3-m
telescope at CTIO.

After applying flat-field corrections, we solved for the astrometry of the
field using the astrometry.net package (Barron et al., 2008; Lang et al.,
2010). We then measured $5^{\prime\prime}$ aperture magnitudes of the
transient and host galaxy using the IRAF apphot package. We obtained archival
$grizy$ magnitudes of several stars in the field with well-defined magnitudes
from Pan-STARRS1 (Chambers et al., 2016; Flewelling et al., 2016). We
calculated SDSS $ugriz$ magnitudes for these stars using the conversions found
in Finkbeiner et al. (2016) and used these to calibrate the transient
magnitudes measured in our follow-up $ugriz$ data. $BV$ follow-up data were
calibrated using stars in SDSS and the $ugriz$ to $BV$ transformations from
Lupton 2005 555http://www.sdss.org/dr5/algorithms/sdssUBVRITransform.html. A
full summary of the photometry, including Swift and ASAS-SN photometry, is
presented in Table 3.

In order to obtain a more accurate position of the transient than that
provided by ASAS-SN, which has $7^{\prime\prime}$ pixels, we solved the
astrometry for two Swope $g$-band images when the transient was bright and
dim. After subtracting these two images to produce an image containing only
the transient, we then calculated the position of the transient using the IRAF
imcentroid package. This yields a position of ($\alpha,\delta$) $=$
(22:43:42.866, $-$16:59:08.410) for ASASSN-18jd. We also found the center of
the host galaxy using the archival $g$-band Pan-STARRS1 image to be
($\alpha,\delta$) = (22:43:42.862, $-$16:59:08.309), yielding an angular
offset of 0$\aas@@fstack{\prime\prime}$12$\pm$0$\aas@@fstack{\prime\prime}$01,
where the uncertainty only incorporates potential uncertainty with the
centroid positions. We used the same procedure to measure the positions of
several stars in both the Swope $g$-band image and the Pan-STARRS $g$-band
image, and find an average random offset of 0$\aas@@fstack{\prime\prime}$20
between the two images. Our measured offset of the transient is thus
consistent with the center of the host, given the random offsets between the
images.

###  2.6 Optical spectroscopy

We obtained multi-epoch optical spectra of ASASSN-18jd spanning 240 d from
2018-05-13 (MJD 58251.4; +34 d) until 2019-06-28 (MJD 58662.4; +445 d) using
the Robert Stobie Spectrograph (RSS, Burgh et al. 2003; Kobulnicky et al.
2003) on the South African Large Telescope (SALT, Buckley et al. 2006), the
Multi-Object Double Spectrographs (MODS, Pogge et al. 2010) on the dual 8.4-m
Large Binocular Telescope (LBT), the Wide Field Reimaging CCD Camera (WFCCD,
Weymann et al. 2001) on the 2.5-m du Pont telescope at Las Campanas
Observatory, the Kast Double Spectrograph on the 3-m Shane Telescope at Lick
Observatory, and the Supernova Integral Field Spectrograph (SNIFS, Lantz et
al. 2004) on the University of Hawaii 88-in telescope. A synopsis of the
spectra is given in Table 5.

The majority of these spectra were reduced using standard IRAF/PyRAF
procedures, including bias subtracting, flat fielding, wavelength fitting
using comparison arc lamps, and flux calibration using spectroscopic standard
stars. The MODS spectra were reduced using the MODS spectroscopic
pipeline666http://www.astronomy.ohio-state.edu/MODS/Software/modsIDL/. Flux
calibration with SALT is difficult because of the telescope design, which has
a moving, field-dependent and under-filled entrance pupil. Observations of
spectrophotometric flux standards can, at best, only provide relative flux
calibration (see, e.g., Buckley et al. 2018), which mostly accounts for the
telescope and instrument sensitivity changes as a function of wavelength.

To combat these issues, we calibrated our SALT spectra using the ground-based
photometry. We re-measured our Swope and Las Cumbres Observatory data using
smaller 1$\aas@@fstack{\prime\prime}$5 apertures to approximate the slit
widths of the spectra and thus account for host contamination. We then
extracted synthetic photometry from our spectra and compared the differences
as a function of the central wavelength of the filters. Fitting a line between
these differences, we scaled the spectra so as to match the smaller aperture
photometry. We repeated this process for the du Pont, Kast, MODS, and SNIFS
spectra so that that our data reduction was consistent. We also corrected for
Galactic extinction. In general, we focus on the emission and absorption
features in the spectra rather than the continuum shape.

The SNIFS spectra were taken from Maunakea, HI at a relatively high airmass,
and some of the spectra were obtained under poor weather conditions. As a
result, the spectra have a very weak blue continuum, and only some of the
spectra show faint emission features blueward of H$\beta$. Additionally,
SNIFS’s dichroic is located at 4800$-$5300 Å (host rest-frame $\sim$
4300$-$4800 Å), making measurements in that range unreliable. For this reason,
we only show the SNIFS spectra redward of rest-frame 4800 Å.

###  2.7 HST/STIS UV spectroscopy

We obtained 6 observations using HST’s Space Telescope Imaging Spectrograph
(STIS; Woodgate et al. 1998) and the FUV/NUV MAMA detectors. We used the
$52\aas@@fstack{\prime\prime}{0}\times 0\aas@@fstack{\prime\prime}{2}$ slit
and the G140L (1150$-$1730 Å, FUV-MAMA) and G230L (1570$-$3180 Å, NUV-MAMA)
gratings. The details of the exposures for each epoch are shown in Table 6.
The source was clearly detected in the two-dimensional frames and spatially
unresolved, so we used the standard HST pipeline for producing one-dimensional
spectra. We performed inverse-variance-weighted combinations of the individual
exposures, merged the FUV and NUV channels, and corrected for Galactic
extinction.

##  3 SED Analysis

Figure 2: Swift UVOT and ground-based photometry of ASASSN-18jd that has been
corrected for Galactic extinction and host-subtracted with the FAST-derived
SED. The magnitudes are in the AB system. Circles correspond to data from
Swift, diamonds from Swope, squares from Las Cumbres, and triangles from
SMARTS. Swift $B$ and $V$ magnitudes were converted to Johnson magnitudes to
allow for comparisons to the ground-based photometry.

We correct our photometry for Galactic extinction and subtract the synthetic
host magnitudes from Table 1 to create UV-optical light curves and temporally-
resolved SEDs of the event. The evolution of the extinction-corrected, host-
subtracted photometry is shown in Figure 2. Our photometry shows that
ASASSN-18jd is fading, albeit slowly. The average decay rates in rest frame
days in the Swift UVW1 and UVW2 filters were $3.3\rm~{}mmag~{}d^{-1}$.

While the transient is fading overall, there are multiple bumps in the light
curve near +90 d, +150 d, and +260 d. These bumps are especially prominent in
the UV, although there appear to be similar, weaker bumps in the $B$\- and
$g$-band light curves. When we remove the roughly linear decay seen in the
light curve, the RMS variability in Swift UVW2, UVM2, UVW1, $U$, and $B$ bands
is 0.15, 0.14, 0.12, 0.12 and 0.13 mag, respectively, all of which are larger
than the median errors of 0.5, 0.5, 0.5, 0.7, and 0.12 mag, respectively. This
is unusual for TDEs, which generally show fairly smooth declines (Holoien et
al., 2019b). For example, the light curve of ASASSN-19bt had RMS variability
of only $\sim$0.01 mag. However, these variations are not unusual for AGNs,
which vary stochastically.

Figure 3: SED of ASASSN-18jd showing the average Swift UVOT data across a
range of epochs as colored circles. The average blackbody fits (see Section 3)
in these epoch ranges are shown as similarly colored lines. The average power-
law fit (see Section 3) is shown as a dotted line. We show the FAST model of
the host galaxy SED as magenta triangles (see Section 2.1) and the time-
averaged optical and UV spectra (see Sections 5 and 6) as grey lines. We show
the lower limits of the H- and He+-ionizing luminosities based on the fluxes
of the broad H$\alpha$ and He ii $\lambda$1640 features (see Sections 5.1 and
6) as red and cyan triangles. We show the mean 0.3$-$10 KeV X-ray luminosity
of the Swift and XMM-Newton spectra (see Section 4) as a black circle. Figure
4: Evolution of ASASSN-18jd as compared to several well-studied TDEs based on
blackbody fits to the Swift UVOT photometry. The events are color-coded as
ASASSN-18jd (black diamonds), ASASSN-14ae (red squares), ASASSN-14li (blue
circles), ASASSN-15oi (yellow triangles), iPTF16fnl (pink pentagons), and
PS18kh (green hexagons). The top-left, top-right, and bottom-left panels show
the luminosity, temperature, and effective radius ($R_{\text{eff}}$)
evolution, respectively. The bottom-right panel compares $R_{\text{eff}}$ to
the tidal radius $R_{\text{T}}$ of the host SMBH of the respective TDEs
assuming a Sun-like star. Except for ASASSN-18jd and PS18kh (Holoien et al.,
2019a), the SMBHs masses were taken from Wevers et al. (2017, 2019a). The
error bar shows the average scatter of $R_{\text{T}}$ caused by uncertainties
in $M_{\text{BH}}$.

We fit blackbody models to the host-subtracted Swift fluxes using Markov chain
Monte Carlo (MCMC) methods. A complete SED of ASASSN-18jd is shown in Figure
3. The median Swift photometry and blackbody fits for four ranges of epochs
are shown. The FAST model of the host SED is added for comparison. We show the
evolution of the bolometric luminosity, effective temperature, and effective
radius in Figure 4. All 6 Swift filters were not available for some epochs,
and the uncertainties on our fits are larger in the epochs with missing
filters. Two epochs had only one or two filters and are not included in the
blackbody fits.

The top-left panel of Figure 4 shows the evolution of the luminosity of
ASASSN-18jd as compared to the TDEs ASASSN-14ae (Holoien et al., 2014; Brown
et al., 2016), ASASSN-14li (Holoien et al., 2016a; Brown et al., 2017),
ASASSN-15oi (Holoien et al., 2016b, 2018), iPTF16fnl (Brown et al., 2018), and
PS18kh (Holoien et al., 2019a). As seen in the photometry, the overall
luminosity of the transient decreases very slowly. We calculate a maximum
luminosity of $L_{\text{max}}=4.5^{+0.6}_{-0.3}\times
10^{44}\rm~{}erg~{}s^{-1}=1.2\times 10^{11}~{}L_{\odot}$. Our mass estimate
for the SMBH yields an Eddington luminosity of $L_{\text{Edd}}=4.9\times
10^{45}\rm~{}erg~{}s^{-1}$, and thus the maximum luminosity corresponds to
$L_{\text{max}}/L_{\text{Edd}}=0.092$. Because of the seasonal gap before
discovery, we likely did not observe the event at peak brightness, and thus
the peak luminosity was probably closer to $L_{\text{Edd}}$. The bumps seen in
the photometry are also seen in the luminosity evolution.

We fit the evolution of the luminosity in rest frame days as a power law
$L\propto(t-t_{0})^{-\alpha}$. Because ASASSN-18jd was discovered after a
seasonal gap and as it was fading, it is possible that the transient started
up to $-$116.4 d before the discovery on MJD 58217.4. This affects the best
fit for $\alpha$. To show this, we compute power-law fits using two different
$t_{0}$, corresponding to the beginning ($t_{1}=58101.0$) and the end
($t_{2}=58217.4$) of the seasonal gap. These yield $\alpha_{1}=0.856\pm 0.049$
($\chi^{2}_{\nu}=3.60$, dof $=$ 46) and $\alpha_{2}=0.431\pm 0.031$
($\chi^{2}_{\nu}=5.56$, dof $=$ 46), respectively. Thus, if the decay follows
a power law, it must have an index of $0.43\lesssim\alpha\lesssim 0.86$, which
is slower than the “canonical” TDE decay rate of $t^{-5/3}$ yet faster than
the $t^{-5/12}$ disc-dominated model of Lodato & Rossi (2011). We also fit the
evolution of the luminosity as an exponential, $L\propto\text{e}^{-t/\tau}$,
where a changing $t_{0}$ does not affect the best-fitting parameters. The best
fit for the exponential profile is $\tau=311.9\pm 14.5\rm~{}d$
($\chi^{2}_{\nu}=2.24$, dof $=$ 46). At early times, this decay rate is much
slower than $t^{-5/3}$. Formally, with such high $\chi^{2}_{\nu}$, all of the
fits are poor, driven by the multiple bumps in the light curve.

The top-right panel of Figure 4 shows the evolution of the effective
temperature of ASASSN-18jd compared to the effective temperatures of the TDEs.
The effective temperature of ASASSN-18jd has remained roughly constant at
$T=(2.5\pm 0.3)\times 10^{4}$ K, which is fairly typical of TDEs. There is
some short-timescale variability in the effective temperature that corresponds
to the bumps in the light curve, which we showed earlier as being stronger at
smaller wavelengths. We also fit a power law,
$f_{\lambda}\propto\lambda^{-\alpha}$, to the host-subtracted Swift fluxes for
each epoch. We find an average index of $\alpha=2.35\pm 0.35$, which is
consistent with the $\alpha\simeq 2.3$ predicted for a standard thin accretion
disc at UV/optical wavelengths (Shakura & Sunyaev, 1973). In most epochs, the
power-law fits are slightly better than the blackbody fits, usually with
$\Delta\chi_{\nu}^{2}\simeq 1$. However, a $T\approx 2.5\times 10^{4}\rm~{}K$
blackbody observed with the Swift filters is nearly indistinguishable from a
power law with the same index, and so our data are still consistent with a
single-temperature blackbody spectrum. Moreover, with the addition of the UV
spectra, we can see the SED break from a power law (see Figure 3) to track the
blackbody SED model.

The bottom-left panel of Figure 4 shows the evolution of the effective radius,
$R_{\text{eff}}=(L/4\pi\sigma T^{4})^{1/2}$, of ASASSN-18jd and the TDEs.
While the geometry is unlikely to be spherical, this should provide a
reasonable estimate of the size of the optically thick, continuum-emitting
region. The evolution of the radius is quite typical of TDEs, showing a
monotonic decline in effective radius, but the evolution is slow compared to
most TDEs. These $R_{\text{eff}}$ are $\sim$100 larger than the Schwarzschild
radius, $R_{\text{Sch}}=1.2^{+1.8}_{-0.7}\times 10^{13}\rm~{}cm$, of the SMBH
given our mass estimate. Finally, the bottom-right panel of Figure 4 compares
the effective blackbody radius and the tidal radius of the SMBH for
ASASSN-18jd and the other TDEs. We use our SMBH mass estimate of
$M_{\text{BH}}=10^{7.6}\rm~{}M_{\odot}$ for ASASSN-18jd and the SMBH masses
from Wevers et al. (2017), Holoien et al. (2019a), and Wevers et al. (2019a)
for the other TDEs, and we compute the tidal radius,
$R_{\text{T}}=R_{\ast}(M_{\text{BH}}/M_{\ast})^{1/3}$, assuming a Sun-like
star ($\rm 1~{}R_{\odot},1~{}M_{\odot}$). For ASASSN-18jd,
$R_{\text{T}}=2.4^{+0.8}_{-0.7}\times 10^{13}\rm~{}cm$. While the effective
radii shown in Figure 4 span nearly two orders of magnitude among the
different TDEs, the ratio of $R_{\text{eff}}/R_{\text{T}}$ spans closer to one
order of magnitude.

Integrating the blackbody luminosity over the span of Swift observations in
rest frame days, we find a total emitted energy of $E=9.6^{+1.1}_{-0.6}\times
10^{51}\rm~{}erg$. This corresponds to an accreted mass of
$M_{\text{acc}}\simeq 0.054\rm~{}M_{\odot}$, for an accretion efficiency of
$\eta=0.1$. This is larger than the accreted mass estimates of most TDEs over
similar timescales (Holoien et al., 2019a; van Velzen et al., 2019b). Recent
TDEs studies have also shown that a significant amount of energy is radiated
prior to peak brightness (Holoien et al., 2019a; Leloudas et al., 2019;
Holoien et al., 2019b), meaning that a significant amount of energy was
probably radiated prior to our first detection.

##  4 X-ray Data

Figure 5: The evolution of the Swift XRT (red squares) and XMM-Newton (blue
squares) 0.3$-$10 keV X-ray luminosities compared to the UV/optical blackbody
luminosity (black diamonds) derived from the blackbody fits. Upper-limits
derived from the Swift XRT data are plotted as red triangles. We show the
ROSAT upper limit (red-dashed line) and the best fitting exponential-decay
model for the blackbody luminosity (black-dashed line). We also plot the
binned flux limit derived from the late-time Swift XRT epochs (MJD > 58500).

In Figure 5, we compare the X-ray luminosity evolution as derived from the
Swift XRT and XMM-Newton observations to the UV/optical luminosity evolution.
Here the X-ray luminosity is estimated from the count rate using
WebPIMMS777https://heasarc.gsfc.nasa.gov/cgi-bin/Tools/w3pimms/w3pimms.pl
assuming a $\Gamma=1.75$ power law as derived from our X-ray spectra and the
line-of-sight Galactic H i column density of $N_{\text{H}}=2.71\times
10^{20}\rm~{}cm^{-2}$ from Kalberla et al. (2005).

The X-ray flux varies by roughly an order of magnitude
($\sim$1042.2$-$$10^{43.2}\rm~{}erg~{}s^{-1}$) including an X-ray flare with a
peak luminosity of $\sim$10${}^{43.7}\rm erg~{}s^{-1}$ around +140 d (MJD
58354.2, ObsID:sw00010680024). By the time of the next Swift XRT observation,
approximately 5 d later, the X-ray emission decreases to previously observed
values. This flare occurs $\sim$10 d prior to a peak in the UV/optical
luminosity near +150 d (see Figures 2, 4, and 5). While the X-ray flare is
quite short, only occurring over a few days, the UV flare near +150 d occurs
on a longer timescale of some tens of days, implying that they are not
necessarily related. Aside from the short flare, the X-ray light curve tends
to follow a similar shallow decline to the optical/UV emission before fading
completely after +275 d. We bin the observations after +275 d to get an upper
limit of $L_{\text{x,late}}<1.1\times 10^{42}\rm~{}erg~{}s^{-1}$.

Figure 6: Evolution of the X-ray Hardness Ratio defined by
$\textit{HR}=(H-S)/(H+S)$ where $H$ is the number of counts in the 2.0$-$10.0
keV energy range while $S$ is the number of counts in the 0.3$-$2.0 keV energy
range. We show the HRs derived from the individual Swift detections (red
circles), the 3$\sigma$ upper-limits (red triangles), and the two deep XMM-
Newton observations (blue squares).

In Figure 6, we show the evolution of the standard hardness ratio HR as a
function of time. ASASSN-18jd shows significant X-ray color evolution with
time, varying between HR $=-0.75\rm~{}and~{}0$. However, during the flare seen
in the X-ray light curve around +140 d, the X-ray emission significantly
softens, before hardening again after the flare ends. As the X-ray emission is
well described by an absorbed blackbody and power-law component (see
discussion below about the X-ray spectra), the softening during the X-ray
flare may be the result of the blackbody component becoming stronger to create
the increase in the soft X-ray flux. Large variations in hardness ratio are
not seen other well known X-ray emitting TDEs, such as ASASSN-14li and
ASASSN-15oi, which tend to show very little variability in their hardness
ratio with time (Auchettl et al., 2018; Holoien et al., 2018). This variation
and softening as the flare peaks is similar to what is seen in AGNs, but in
AGNs this softening is due to a steepening of the power-law component, whereas
the blackbody component remains roughly constant (see Figure 4 from Auchettl
et al. 2018). This is different from what we see in ASASSN-18jd, where the
power-law index remains fixed and the blackbody component fluctuates.

Observation | Epoch [MJD] | $\Gamma$ |  $kT$ [keV] |  $L_{\text{pwl}}$ [$10^{43}$ erg s-1] |  $L_{\text{BB}}$ [$10^{43}$ erg s-1] |  $R_{\text{BB}}$ [$10^{10}$ cm]
---|---|---|---|---|---|---
XMM-Newton 1 | 58266.8 | $1.7\pm 0.1$ | $0.08\pm 0.02$ | $1.1\pm 0.5$ | $0.31\pm 0.07$ | $9.0^{+34.8}_{-5.7}$
XMM-Newton 2 | 58280.6 | $1.7\pm 0.2$ | $0.14\pm 0.04$ | $1.4\pm 0.7$ | $0.37\pm 0.06$ | $2.5^{+2.4}_{-1.0}$
Binned Swift | 58249.8$-$58390.5 | $1.7_{-0.3}^{+0.4}$ | $0.07\pm 0.01$ | $0.54\pm 0.01$ | $1.2\pm 0.3$ | $31^{+19}_{-11}$
Table 2: Parameters derived for the power-law and blackbody components of the
three X-ray spectra. Figure 7: The two XMM-Newton spectra and the binned Swift
XRT spectrum. For the XMM-Newton spectra, the red and blue data are from the
two MOS cameras, and the green data are from the PN camera. The top panels
shows the data (crosses), the blackbody and power-law model fits (dashed
lines), and the combined model fits (solid line). The bottom panels are the
ratios between the model and the data.

We extract spectra from our two deep XMM-Newton observations and the merged
Swift observations spanning the first and last $\sim$200 d of our study. Due
to the lack of X-ray emission at late times (see Figure 5), our second merged
Swift spectrum does not have enough signal to characterize the emission and so
we do not consider it further. In Figure 7, we show the unconvolved PN and MOS
spectra from our two XMM-Newton observations and our merged Swift XRT spectrum
with their best fit models and residuals. We find that all three spectra are
well fit by an absorbed blackbody plus power-law model. Initially, the column
density ($N_{\text{H}}$), blackbody temperature ($kT$), power-law index
($\Gamma$) and normalizations of each model were free parameters, but we find
that $N_{H}$ is unconstrained, so we fix it to the line-of-sight value. We
summarize the parameters derived from the X-ray spectra in Table 2. Fitting
each spectrum with only one of the two components produces a significantly
worse fit. Within the uncertainties, we find no evidence for changes in the
power-law indices or blackbody temperatures between observations. The derived
X-ray blackbody temperatures are consistent with those found for other X-ray
bright TDEs, such as ASASSN-14li (Brown et al., 2017) and ASASSN-15oi (Gezari
et al., 2017a; Holoien et al., 2018). and are on the low-temperature tail of
blackbody temperatures found for AGNs (Reynolds, 1997; Ricci et al., 2017).
The observed power-law index is seen in both TDEs such as ASASSN-15oi (Gezari
et al., 2017a; Holoien et al., 2018) and Swift J1644+57 (Burrows et al., 2011;
Bloom et al., 2011) and in AGNs (Ricci et al., 2017). The mean of the total
0.3$-$10 keV luminosity derived from the three X-ray spectra is shown in
Figure 3.

The luminosities of the blackbody and power-law components found for the two
XMM-Newton observations are consistent given the uncertainties. However, our
merged Swift spectrum for the first 200 d of emission, which includes the
flare near +140 d, has a blackbody component that is approximately an order of
magnitude more luminous, while the power-law component is slightly less
luminous. The increase in the blackbody flux likely corresponds to an increase
in the X-ray effective radius, as $kT$ remains roughly constant. The apparent
expansion of the X-ray emitting region during the flare is very different from
the shrinking radius of the UV emitting region (see Figure 4).

The X-ray-derived blackbody radii of $3\times 10^{10}$ to $3\times
10^{11}\rm~{}cm$ are smaller than the Schwarzschild radius of the SMBH by a
factor of roughly 30 to 300. These radii are also smaller than those found for
the X-ray blackbodies of the TDEs ASASSN-14li and ASASSN-15oi (Brown et al.,
2017; Holoien et al., 2018), which were both of order 10${}^{12}\rm~{}cm$.
Additionally, the X-ray blackbody radii of ASASSN-14li and ASASSN-15oi were
1$-$10 times larger than the associated SMBH Schwarzschild radii, rather than
two orders of magnitude smaller. The X-ray blackbody radii for ASASSN-18jd
are, however, similar to that of PS18kh, which also showed X-ray blackbody
radii smaller than the predicted Schwarzschild radius by a factor of $\sim$100
(van Velzen et al., 2019a). In AGNs, the blackbody component (also called the
“soft-excess”) is thought to be created through other physical processes like
Comptonization and reprocessing, and thus a blackbody radius is not a
meaningful quantity.

##  5 Optical spectroscopy

Figure 8: The du Pont (blue), SALT (black), Kast (magenta), and MODS (red)
spectra. Prominent spectral features are labelled, and regions strongly
affected by telluric absorption are shaded. The du Pont, SALT, Kast, and MODS
spectra were smoothed with 10, 5, 7, and 7 Å-wide bins, respectively. The gaps
in the SALT spectra are due to the spectrograph’s design. Figure 9: Multi-
epoch SNIFS spectra of ASASSN-18jd. Regions strongly affected by telluric
absorption are shaded.

The SALT, du Pont, Kast, and MODS spectra are presented in Figure 8. Because
of the smaller wavelength coverage of the SNIFS spectra, they are presented
separately in Figure 9. In both figures, prominent spectral features are
labelled, and telluric bands are shaded. The most prominent features in the
spectra are strong H$\alpha$, H$\beta$, and H$\gamma$ lines and a blue
continuum. There are multiple features around 4600 Å and a broad (FWHM $\sim
100$ Å) feature centered near 3430 Å. There are also collisionally excited
lines like [O ii] $\lambda$3727, [O iii] $\lambda\lambda$4959,5007, [N ii]
$\lambda\lambda$6548,83, and [S ii] $\lambda\lambda$6716,31 that are commonly
seen in galaxy spectra. We can also see Ca ii $\lambda\lambda$3934,68 and Mg i
$\lambda\lambda$5173,84 absorption from the host and Na i
$\lambda\lambda$5890,96 absorption from both our Galaxy and the host.

Figure 10: Continuum-subtracted H$\alpha$ and H$\beta$ profiles from the
optical spectra. For the H$\alpha$ profiles, the black lines show the spectra,
and the blue, yellow, and red lines show the best fitting Gaussian profiles
for the narrow H$\alpha$, [N ii] $\lambda\lambda$6548,83, and broad H$\alpha$
components, respectively. For the H$\beta$ profiles, the black lines show the
spectra, and the magenta and cyan lines show the best fitting Gaussian
profiles for the narrow and broad H$\beta$ components, respectively. The
spectra were scaled to one another by holding the integrated flux of the
narrow H$\alpha$ emission constant. The spectra are labelled by date and
source.

In Figure 10, we show the H$\alpha$ and H$\beta$ emission profiles. We
decompose the H$\alpha$ profile into a sum of three narrow Gaussian profiles
centered on H$\alpha$ and [N ii] $\lambda\lambda$6548,83 and one broad
Gaussian profile centered on H$\alpha$. We repeat this process for the
H$\beta$ profiles in the du Pont, Kast, and MODS spectra (see Figure 8), using
only one broad and one narrow Gaussian profile centered on H$\beta$. Because
of the relatively low spectral resolution ($R\sim 300$) of the SALT spectra
(see Figure 8), we are unable to decompose the Balmer line profiles.
Similarly, we are unable to decompose the H$\beta$ profiles in the SNIFS
spectra due to noise (see Figure 9). The H$\gamma$ feature is different from
the other Balmer features, showing a broad, flat-topped profile. Such
H$\gamma$ profiles are sometimes seen in the spectra of broad-line AGNs. The
most likely explanation for the irregular profile is contamination from nearby
[O iii] $\lambda$4363, so we do not try to decompose the H$\gamma$ profile
into narrow and broad components.

The narrow-line emission does not appear to change over time, and so it is
likely coming from the host galaxy. Assuming the narrow-line emission is
unrelated to the transient, we use the flux of the narrow H$\alpha$ component
as a constant to scale the spectra across multiple epochs. We also measure the
ratios of the narrow-line emission in the spectra, finding
$\rm\log{([\text{N\,{ii}}]~{}\lambda 6583/H\alpha)}=-0.3\pm 0.1$,
$\rm\log{([\text{O\,{iii}}]~{}\lambda 5007/H\beta)}=0.0\pm 0.1$,
$\rm\log{([\text{S\,{ii}}]~{}\lambda\lambda 6716,32/H\alpha)}=-0.7\pm 0.1$.
These ratios correspond to the “composite” region between the AGN-dominated
and SF-dominated regions of the BPT diagrams (Kewley et al., 2001; Kauffmann
et al., 2003). This fits with our understanding of the host as having only
weak, if any, AGN activity before the transient. We measure the narrow-line
H$\alpha$ emission to be $F_{\text{H$\alpha$,n}}=(6.0\pm 0.7)\times
10^{-15}\rm~{}erg~{}s^{-1}~{}cm^{-2}$, corresponding to a luminosity
$L_{\text{H$\alpha$,n}}=2.2\times 10^{41}\rm~{}erg~{}s^{-1}$. We also measure
the [O ii] $\lambda$3727 emission to be $F_{\text{[\text{O\,{ii}}]}}=(1.0\pm
0.5)\times 10^{-15}\rm~{}erg~{}s^{-1}~{}cm^{-2}$, corresponding to a
luminosity $L_{\text{[\text{O\,{ii}}]}}=3.7\times 10^{40}\rm~{}erg~{}s^{-1}$.
These luminosities correspond to SFR rates of
SFR${}_{\text{H$\alpha$}}=(1.8\pm 0.2)\rm~{}M_{\odot}~{}yr^{-1}$ and
SFR${}_{[\text{O\,{ii}}]}=(0.5\pm 0.2)\rm~{}M_{\odot}~{}yr^{-1}$ (Kennicutt,
1998). These SFR rates are reasonably consistent with the estimate of SFR
$=0.6^{+0.1}_{-0.3}\rm~{}M_{\odot}~{}yr^{-1}$ found from the FAST SED models
in Section 2.1, suggesting that much of the narrow-line emission is due to
star formation.

###  5.1 Balmer features

Figure 11: Top: Evolution of the broad H$\alpha$ and H$\beta$ flux over time.
Error-bars are 1$\sigma$. The H$\beta$ fluxes are scaled by a factor of 2.8,
the expected ratio of H$\alpha$/H$\beta$ for recombination. Bottom: Dependence
of the flux on the line FWHM for the broad H$\alpha$ and H$\beta$ profiles.
The SNIFS, du Pont, Kast, and MODS spectra are represented by circles, boxes,
diamonds, and triangles, respectively.

The evolution of the flux and FWHM of the broad H$\alpha$ and H$\beta$
features are shown in Figure 11. To estimate the uncertainties, we resample
the spectra with bootstrapping methods. In each trial a spectral data point is
randomly sampled $n_{i}$ times. For $n_{i}>0$, we reduce the error to
$\sigma_{i}/\sqrt{n_{i}}$, while for $n_{i}=0$, we double the error. We do
this 100 times and use the mean and dispersion of the results for Figure 11.
After excluding the points with the largest error bars, the broad H$\alpha$
features have an average FWHM of $3250\pm 350\rm~{}km~{}s^{-1}$. In the du
Pont spectra, the broad H$\beta$ features, which are relatively noisy, have
FWHM of $5900\pm 1000\rm~{}km~{}s^{-1}$ and $5900\pm 400\rm~{}km~{}s^{-1}$. In
the Kast spectra, the H$\beta$ features have similar FWHMs of $3600\pm
400\rm~{}km~{}s^{-1}$ and $4100\pm 300\rm~{}km~{}s^{-1}$. The MODS spectra
yield very different H$\beta$ FWHMs of $3200\pm 200\rm~{}km~{}s^{-1}$ and
$7000\pm 200\rm~{}km~{}s^{-1}$. This last measurement is peculiarly large,
especially considering the fact that the H$\alpha$ FWHM is not anywhere near
as large. While Figure 11 may appear to show an inverse correlation between
the flux and FWHM of the broad H$\alpha$ feature, the three points with larger
flux and smaller FWHM values are also from the noisiest spectra. Additionally,
the H$\beta$ feature appears to show the opposite trend. The most likely
explanation is that these features are roughly constant over time, and the
fluctuations seen in Figure 11 are dominated by noise.

Generally, the Balmer features of H-rich TDEs fade and become narrower as the
continuum fades. ASASSN-14li, for example, showed a roughly linear correlation
between the continuum and the H$\alpha$ luminosities and line width while
fading (Holoien et al., 2016a; Brown et al., 2017). By contrast, the Balmer
features of AGNs become fainter ($\propto L^{0.5}$) but broader ($\propto
L^{-0.25}$) as the continuum luminosity ($L$) fades (Korista & Goad, 2004;
Peterson et al., 2004; Denney et al., 2009). ASASSN-18jd shows neither trend
despite the continuum having faded by a factor 3$-$5 by the time of the last
epoch where spectra were taken.

We use the flux of the broad component of the H$\alpha$ emission to estimate
the unobserved H-ionizing luminosity. The average broad H$\alpha$ line flux of
$F_{\text{H$\alpha$,b}}=(2.5\pm 0.4)\times
10^{-14}\rm~{}erg~{}s^{-1}~{}cm^{-2}$ corresponds to an H$\alpha$ luminosity
of $L_{\text{H$\alpha$,b}}=9.4\times 10^{41}\rm~{}erg~{}s^{-1}$, which implies
a minimum H-ionizing luminosity of $L_{\text{H\,{ii}}}=1.6\times
10^{43}\rm~{}erg~{}s^{-1}$, assuming Case B recombination. While this is
comparable to the H-ionizing luminosity of other H-rich TDEs (Brown et al.,
2017), it is several orders of magnitude smaller than our predicted blackbody
luminosity at the 912 Å Lyman continuum break from the UV/optical blackbody
fits shown in Figure 3. This implies a small covering fraction for the line
emitting gas, which is also true in the BLRs of AGNs. If we assume the
$R_{\text{BLR}}$$-$$L$ relationship from Bentz et al. (2009) holds, then the
BLR radius has a range of $R_{\text{BLR}}\sim(5\pm 2)\times 10^{14}\rm~{}cm$,
which places the material outside the tidal disruption radius for a Sun-like
star, yet inside the photospheric radius estimated from the blackbody fits.
Similarly, if we interpret the line widths as an estimate of the escape
velocity using our average H$\alpha$ width of FWHM $=3250\rm~{}km~{}s^{-1}$,
we get a radius for the Balmer emission of $R_{\text{H$\alpha$,esc}}\sim
1\times 10^{17}\rm~{}cm$, which is outside the photospheric radius.

###  5.2 He and metal lines

Figure 12: Multi-epoch SALT spectra of the emission features near 4600 Å.
Black lines denote He and potential N iii-v lines present in the spectra,
whereas grey lines mark potential Fe ii lines that are observed in AGNs and
some TDEs (Wevers et al., 2019b). Fe ii lines are taken from Kovačević et al.
(2010). Over time, the He ii $\lambda$4686 emission appears to get weaker with
respect to the bluer N iii emission.

The series of features around 4600 Å show multiple peaks at roughly 4510,
4570, 4635, and 4686 Å. These features are visible in Figure 8 and highlighted
in Figure 12. The last of these features is easily attributed to He ii
$\lambda$4686, which is a common feature in TDE spectra (e.g., Arcavi et al.
2014). The other features (and/or features at similar wavelengths) are also
seen in some TDE spectra (see, e.g., Leloudas et al. 2019). A possible origin
of these features is a blend of highly ionized N lines that are seen in Wolf-
Rayet (WR) stars (see, e.g., Crowther 2007), in particular N iii
$\lambda$4510, N v $\lambda$4610, and N iii $\lambda\lambda$4634,40.
Alternatively, this emission could be due to Fe ii, which is commonly seen in
AGN spectra and has been recently identified in the spectra of the TDEs
ASASSN-15oi and ASASSN-18ul (Wevers et al., 2019b). Commonly-used Fe ii
templates (e.g., Kovačević et al. 2010) list the brightest lines in this
wavelength regime as $\lambda$4418, $\lambda$4523, $\lambda$4549,
$\lambda$4583, and $\lambda$4629, where the last line is actually a blend of
two lines whose relative strengths vary. N iii-v seems the more likely
contributors to the emission profile, although some features line up with Fe
ii lines, and there are no strong N iii-v lines near the feature at 4570Å, so
N iii-v cannot be the sole source of emission in this wavelength region.

In addition to the features around 4600 Å, we can see a prominent emission
feature near H$\delta~{}\lambda 4101$ in the +34 d, +90 d, and +97 d spectra.
In the +34 d spectrum, this feature appears stronger than H$\gamma$, implying
that the emission at 4100 Å must be contaminated by emission from another
source. The most likely line is N iii $\lambda$4100, a commonly observed line
in WR spectra that is of equal strength to N iii $\lambda$4640 and that has
been identified in other TDE spectra (Blagorodnova et al., 2019; Leloudas et
al., 2019).

As we see in Figure 12, the He ii $\lambda$4686 line appears to get weaker
over time compared to the nearby N iii $\lambda$4640\. At earlier times, the
He ii emission is nearly as strong as the N iii, whereas at later times, the
He ii line is only barely visible. For comparison, ASASSN-14li and ASASSN-18pg
showed the opposite behavior, where the N iii line started brighter and became
fainter than the He ii line (Holoien et al., 2016a; Leloudas et al., 2019).
Additionally, the H$\delta~{}\lambda$4101 + N iii $\lambda$4100 emission
feature appears quite prominent and quite broad (FWHM $\sim$ 100 Å) in the +34
d spectrum, and it appears fainter in the +90 d and +97 d spectra. In the +210
d spectrum, the emission is both fainter and narrower, and in the +421 d and
the +445 d spectra, the feature does not appear at all (see Figure 8).

Throughout the epochs, we can see a broad featured centered near 3430 Å. In
the +210 d MODS spectrum, we can see a slightly brighter yet narrower bump
centered near 3130 Å. Recent studies of TDE spectra (Blagorodnova et al.,
2019; Leloudas et al., 2019) have identified similar features at these
wavelengths as O iii $\lambda$3133 and O iii $\lambda\lambda$3428,44. These
studies have also proposed that these O iii emission lines, as well as the N
iii emission lines at N iii $\lambda$4100 and N iii $\lambda$4640, are due to
Bowen Fluorescence (BF) with He ii $\lambda$303.78 emission (Bowen, 1934).
However, there are some problems with this interpretation. BF would not
produce any N iii lines in this region besides N iii $\lambda\lambda$4097,4103
and N iii $\lambda\lambda$4634,40 (McClintock et al., 1975; Netzer et al.,
1985; Selvelli et al., 2007). This is in conflict with our identification of N
iii $\lambda$4510\. Additionally, O iii $\lambda$3133 should be roughly 3
times stronger than O iii $\lambda$3444 (McClintock et al., 1975; Netzer et
al., 1985; Liu et al., 1993; Kastner & Bhatia, 1996; Selvelli et al., 2007).
This is not the case in our spectra, though the exact strengths are difficult
to estimate because these are broad features and could be blended with other
lines, such as [Ne v] $\lambda$3426\. However, an additional problem arises
from the fact that in the +445 d MODS spectrum, the O iii $\lambda$3133
feature is barely visible if at all, whereas the O iii $\lambda$3444 feature
is still quite prominent. These inconsistencies are difficult to explain with
BF, so it may be that BF is not responsible for the emission at 3430 Å and
3430 Å or for the N iii emission in ASASSN-18jd.

We observe an emission line centered at 6375 Å that is best seen in the early-
time optical spectra. This may be [Fe x] $\lambda$6375, a coronal (high-
ionization, forbidden) line often seen in AGN spectra. This identification may
be problematic because of its closeness to the telluric band at 7186 Å (host
rest-frame $\sim 6420$ Å). However, this feature also appears to fade over
time, providing evidence that the feature is in fact real. There also appears
to be a feature around 5300 Å that might be due to [Fe xiv] $\lambda$5303,
another coronal emission line seen in AGNs, but this region is affected by
Milky Way absorption from Na i $\lambda\lambda$5890,96 (host rest-frame $\sim$
5260 Å). While these lines are seen in AGNs, notably Narrow Line Seyfert 1s
(NLS1s), these lines are usually accompanied by lower-level ionization lines
like [Fe vii] $\lambda$6088 and are much weaker than emission from [O iii]
$\lambda$5007\. When these lines are brightest in ASASSN-18jd’s spectra, [Fe
vii] $\lambda$6088 is not present, and [Fe x] $\lambda$6375 is roughly
equivalent in strength to [O iii] $\lambda$5007, implying that these lines are
not attributable to “normal” AGN or NLS1 activity. These lines are indicative
of strong soft X-ray flux, as Fe x and Fe xiv have ionization potentials of
234 eV and 361 eV, respectively. The transient nature of the lines, especially
[Fe x] $\lambda$6375, implies that they are associated with the transient
X-ray flux of ASASSN-18jd. Other studies have found similar, transient coronal
Fe lines in galaxies with faint or no AGN activity (Komossa et al., 2009; Wang
et al., 2011; Wang et al., 2012). These galaxies were dubbed Extreme Coronal
Line Emitters (ECLEs) (Wang et al., 2012), and it was proposed that the soft
X-ray flux required to generate these lines could originate from TDEs.

##  6 UV spectroscopy

Figure 13: Multi-epoch UV spectra of ASASSN-18jd from HST/STIS. The spectra
were extinction corrected and smoothed with 1 Å-wide bins. Prominent spectral
features are labelled. All absorption features are due to telluric and MW
absorption and not associated with the transient. The Swift fluxes (red
squares) that were observed in roughly concurrent epochs are scaled in the
plot to match the flux values of the spectra at the filters’ center
wavelengths. Also plotted are the (scaled) blackbody models (dashed-red lines)
derived from the Swift fluxes. The absolute values differ by 20$-$25 per cent.

We present the six epochs of HST/STIS UV spectra of ASASSN-18jd in Figure 13.
There is a $\sim$20$-$25 per cent difference between STIS and Swift fluxes for
concurrent epochs, which is likely explained by slit losses and the large
widths of the Swift filters. We show the scaled Swift fluxes and fitted
blackbody continua in Figure 13 along with the STIS spectra to show the
strengths of line features relative to the continuum. The most prominent
emission features seen for all the epochs are Ly$\alpha$, N v $\lambda$1240,
Si iv $\lambda\lambda$1394,1403, C iv $\lambda$1550, He ii $\lambda$1640, and
C iii] $\lambda$1909, as well as faint bumps attributable to O i $\lambda$1302
and N iii] $\lambda$1302\. There are absorption features at local-rest-frame
Ly$\alpha$, Si ii $\lambda$1260, O i $\lambda$1302/Si ii $\lambda$1304, C ii
$\lambda$1335, Si ii $\lambda$1527, C iv $\lambda$1550, and Mg ii
$\lambda$2800\. There also appears to be a double-peaked feature around 2800 Å
that is most likely Mg ii $\lambda$2800, though it is difficult to be certain
because the feature is close to the detector’s edge.

The lack of strong N iii-v features is in conflict with our identification of
the features at 4100 Å and near 4600 Å as N iii-v emission. Most notable is
the absence of N iv] $\lambda$1486\. Since the critical density of N iv]
$\lambda$1486 is of the same order of magnitude as the critical density for C
iii] $\lambda$1909, with $n_{\text{crit}}\sim
10^{9}$$-$$10^{10}\rm~{}cm^{-3}$, the lack of N iv] $\lambda$1486 is likely
not a consequence of having a particular density for the line-emitting gas.
This lack of N iii-v emission is not an issue if the optical N iii lines are
due to BF, which does not produce N iii-v lines in the NUV or FUV.

Figure 14: Continuum-subtracted Ly$\alpha$ profiles from the STIS spectra. The
black lines are the 1 Å-bin smoothed spectra, and the blue and red lines show
the best fitting Gaussian profiles for the narrow Ly$\alpha$ and broad
Ly$\alpha$ components, respectively. The red-dashed line is the sum of the two
Gaussians. The N v $\lambda$1240 emission is also labelled.

We trace the evolution of the Ly$\alpha$ feature across epochs using the same
method that we use for fitting the H$\alpha$ profiles. We find that the
feature can be roughly decomposed into narrow (FWHM $\sim
1000\rm~{}km~{}s^{-1}$) and broad (FWHM $\sim 7000\rm~{}km~{}s^{-1}$)
components. Additionally, we find that the broad component becomes brighter
during the last epoch of observation. This is true whether or not we hold the
narrow component fixed in flux. The continuum-subtracted Ly$\alpha$ features
and the fits associated with them are presented in Figure 14. Because
Ly$\alpha$ is very sensitive to changes in opacity and optical depth, the
change in this line profile over time implies some dynamical changes in the
system over time which are not discernible from the Balmer lines. The maximum
broad Ly$\alpha$ flux taken from the last epoch of
$F_{\text{Ly$\alpha$}}=(1.7\pm 0.1)\times
10^{-13}\rm~{}erg~{}s^{-1}~{}cm^{-2}$, whereas for the other five epochs, the
average flux is $F_{\text{Ly$\alpha$}}=(8.4\pm 2.3)\times
10^{-14}\rm~{}erg~{}s^{-1}~{}cm^{-2}$.

The broad He ii $\lambda$1640 feature can be used to estimate the He+-ionizing
luminosity. We get an integrated flux of $F_{\text{1640~{}\AA}}\simeq(5\pm
2)\times 10^{-14}\rm~{}erg~{}s^{-1}~{}cm^{-2}$, corresponding to a luminosity
$L_{\text{1640~{}\AA}}=2.5\times 10^{42}\rm~{}erg~{}s^{-1}$. Assuming Case B
recombination and $T=2\times 10^{4}\rm~{}K$ gas, this luminosity implies a
minimum He+-ionizing luminosity of $L_{\text{He\,{iii}}}=2.8\times
10^{43}\rm~{}erg~{}s^{-1}$. We compare this estimate of the He+-ionizing
luminosity to the UV/optical blackbody luminosity in Figure 3. This luminosity
is the same order of magnitude as the H-ionizing flux, yet it is significantly
higher than the predicted blackbody luminosity at the He+-ionizing edge at 227
Å. This implies that there is a significant amount of EUV photons that are
being generated by a currently-unseen radiation source.

Aside from the changes in Ly$\alpha$ and some apparent fluctuations of N v
$\lambda$1240, the UV spectra of ASASSN-18jd do not appear to change with
time. This is in contrast to the optical spectra, where He ii $\lambda$4686, N
iii $\lambda$4100 and [Fe x] $\lambda$6375 weaken over time, albeit over
longer timescales than the differences between the HST epochs.

##  7 Discussion

Figure 15: Left: Early-time optical spectra of ASASSN-18jd compared to known
TDEs. The spectra are not host-subtracted. Right: Early-time optical spectra
compared to a composite quasar spectrum, the TDE candidate PS16dtm, the rapid
turn-on event ZTF18aajupnt, and a N-type WR star. We approximated and
subtracted the continuum of each spectra in order to compare spectral emission
features, which are labelled. Figure 16: UV spectrum of ASASSN-18jd compared
to the UV spectra of known TDEs, a composite quasar spectrum, and a “N-rich”
quasar spectrum. The Swift-derived blackbody continuum fits for ASASSN-18jd
and the TDEs are shown as dotted, colored lines. Prominent spectral features
are labelled.

ASASSN-18jd is a unique event among transients and shares properties with
optically-detected TDEs and with AGNs. We compare the optical and UV spectra
of ASASSN-18jd to well-studied TDEs, other extragalactic transients, and
objects with similar spectra in Figures 15 and 16. In Figure 15, we compare
two early-time optical spectra of ASASSN-18jd to known TDEs (Holoien et al.,
2014, 2016a, 2016b; Hung et al., 2017; Blagorodnova et al., 2019; Hung et al.,
2019), as well as the composite quasar spectrum from Vanden Berk et al.
(2001), an early-time spectrum of the TDE candidate PS16dtm (Blanchard et al.,
2017), an early-time spectrum of the rapid turn-on event ZTF18aajupnt/AT
2018dyk (Frederick et al., 2019), and the spectrum of a WR star with strong He
ii and N iii-v features (Yaron & Gal-Yam, 2012). In Figure 16, we compare a UV
spectrum of ASASSN-18jd to TDEs with published UV spectra (Cenko et al., 2016;
Brown et al., 2018; Blagorodnova et al., 2019; Hung et al., 2019), the
composite quasar spectrum from Vanden Berk et al. (2001), and an “N-rich”
quasar spectrum (Batra & Baldwin, 2014).

###  7.1 ASASSN-18jd as a TDE

The SED of ASASSN-18jd is TDE-like, with a strong blue continuum that is
reasonably-well fit by a blackbody with temperature $T\simeq 2.5\times
10^{4}\rm~{}K$. With a maximum observed luminosity of
$L_{\text{max}}=4.5^{+0.6}_{-0.3}\times 10^{44}\rm~{}erg~{}s^{-1}$,
ASASSN-18jd would be one of the most luminous optically-discovered TDEs to
date, and its host galaxy would be one of the most massive TDE host galaxies.
Our SMBH mass estimate of $\log{M_{\text{BH}}/\rm M_{\odot}}=7.6\pm 0.4$ is
one of the largest masses for a SMBH associated with an observed TDE (Wevers
et al., 2017, 2019a).

Empirically, it has been suggested that the UV/optical fading timescale
increases with SMBH mass (Blagorodnova et al., 2017; Wevers et al., 2017; van
Velzen et al., 2019b). ASASSN-18jd has a decay rate of
$3.3\rm~{}mmag~{}d^{-1}$ in the Swift UVW2 and UVW1 filters. In these terms,
the most similar TDE is GALEX D3-13, which has a SMBH mass estimate of
$\log{M_{\text{BH}}/\rm M_{\odot}}=7.36\pm 0.44$ and a decay rate of $2.6\pm
0.2\rm~{}mmag~{}d^{-1}$ (Wevers et al., 2017). However, GALEX D3-13 has a very
sparsely-sampled light curve compared to ASASSN-18jd, so it is unclear how
well the two decay rates match.

The luminosity evolution of ASASSN-18jd is different from “normal” TDEs in
that it does not smoothly decline. There appear to be multiple bumps in the
light curve where the luminosity returns to near maximum values, similar to
that seen from highly variable AGNs. By comparison, most TDEs with well-
sampled UV/optical light curves show only smooth, monotonic fading. Some
recently discovered TDEs have shown deviations from monotonic fading:
ASASSN-18ul showed a light curve plateau $\sim$40 d after peak brightness
(Wevers et al., 2019b) and ASASSN-19bt showed a short flare before reaching
peak brightness (Holoien et al., 2019b). In a more extreme example, PS18kh
showed a UV re-brightening $\sim$50 d after peak (Holoien et al., 2019a) and
then, after a seasonal gap due Sun constraints, appeared to return to near-
peak luminosity (van Velzen et al., 2019a). However, these deviations from
monotonic fading were still “smooth”, with little short-timescale variation,
which is qualitatively different from the bumpiness in the light curve of
ASASSN-18jd.

As seen in Figure 15, the optical spectra ASASSN-18jd is different from the
spectra of other TDEs. The Balmer features of ASASSN-18jd are narrower than
those in TDEs like ASASSN-14ae, iPTF16axa, and iPTF16fnl (Holoien et al.,
2014; Hung et al., 2017; Blagorodnova et al., 2017). The widths of the Balmer
features are similar to those of ASASSN-14li, which had a maximum H$\alpha$
FWHM of $\sim$3000 km s-1 (Holoien et al., 2016a). However, the Balmer
features of ASASSN-14li showed a clear evolution, becoming dimmer and narrower
as the continuum faded (Holoien et al., 2016a; Brown et al., 2017), whereas
the Balmer features of ASASSN-18jd do not show any coherent evolution.
ASASSN-14li was also discovered after a seasonal gap in ASAS-SN coverage and
as it was fading, and thus its emission lines were probably broader than
$\sim$3000 km s-1 closer to its actual peak brightness.

With the exception of ASASSN-14li, the H-ionizing luminosity estimated from
the H$\alpha$ flux for other TDEs is usually greater than the predicted
blackbody luminosity. This could be related to the lack of evolution in the
Balmer features of ASASSN-18jd. Generally, in both TDEs and reverberation
mapped AGNs, the Balmer line fluxes drop with the continuum luminosity $L$.
For TDEs, this is usually $L_{\text{H$\alpha$}}\propto L$ (Brown et al.,
2017), whereas for AGNs, this is usually $L_{\text{H$\alpha$}}\propto L^{0.5}$
(Korista & Goad, 2004). This is not the case for ASASSN-18jd, where the Balmer
line fluxes remained roughly constant while the continuum luminosity dropped.
A possible explanation is that the broad line emission could be limited by the
amount of broad line gas rather than the luminosity needed to ionize it. In
this “matter bounded” scenario, the line luminosity becomes insensitive to
changes in the ionizing continuum luminosity, which would also explain the
lack of evolution in the FWHM of the Balmer lines. Regardless, the relatively
weak dependence of Balmer flux to continuum seen in ASASSN-18jd is more
similar to AGN activity than emission from TDEs.

The optical spectra and SED evolution of ASASSN-18jd resembles that of the
transient PS16dtm, claimed to be a TDE in a known NLS1. PS16dtm brightened
rapidly ($\sim$50 d), faded slowly, and exhibited a transient blue continuum
and variable, strong Fe ii line emission. While PS16dtm showed an increase in
the UV and optical flux, the upper-limit set by the non-detection of X-ray
flux was an order of magnitude below the pre-flare detections of the host NLS1
by XMM-Newton from years earlier. The light curve of PS16dtm is similar to
that of ASASSN-18jd, in that it faded slowly and also exhibited some bumps
inconsistent with monotonic decay. ASASSN-18jd does not, however, show the
strong Fe ii emission that dominates the spectra of PS16dtm.

As seen in Figure 16, the UV spectrum of ASASSN-18jd is unique compared to
that of TDEs. Whereas TDEs do not show lower ionization level lines like O i
or Mg ii (Cenko et al., 2016; Brown et al., 2018; Blagorodnova et al., 2019),
these lines are very common in AGN spectra. Our UV spectra show weak N v
$\lambda$1240 and N iii] $\lambda$1750, and no N iv] $\lambda$1486 emission.
TDEs observed in the UV like ASASSN-14li (Cenko et al., 2016) and iPTF16fnl
(Brown et al., 2018) showed strong, broad emission from all of these N lines,
especially N v $\lambda$1240\. These emission lines are not seen in iPTF15af
(Blagorodnova et al., 2019), which instead showed strong absorption bands like
those seen in broad absorption line quasars. UV spectra of PS18kh appear to
show strong, somewhat variable N v $\lambda$1240 emission, but no N iii]
$\lambda$1750 or N iv] $\lambda$1486 emission (Hung et al., 2019).
Additionally, our UV spectra show strong C iv $\lambda\lambda$1548,51 and C
iii] $\lambda$1909, the latter of which was not detected in the other four
TDEs. Both of these C iii-iv emission lines are common in AGN spectra, while
strong N iii-iv lines are only seen in the rare N-rich AGNs (e.g., Batra &
Baldwin 2014).

The X-ray light curve of ASASSN-18jd is also unique compared to other X-ray
bright TDEs. ASASSN-14li (Holoien et al., 2016a; Brown et al., 2017), Swift
J1644+77 (Burrows et al., 2005; Bloom et al., 2011; Levan et al., 2016), and
Swift J2058+05 (Cenko et al., 2012) all exhibited X-ray light curves which
decay monotonically and coherently following approximately the canonical
$t^{-5/3}$ at early times, irrespective of any short timescale variability
(Auchettl et al., 2017, 2018). As ASASSN-14li evolved, its light curve was
more consistent with disc emission ($t^{-5/12}$) rather than fall back, but
the overall the decay was still monotonic (Auchettl et al., 2017). While the
UV/optical emission of ASASSN-15oi showed a steady decay, its X-ray emission
did not (Holoien et al., 2016b; Gezari et al., 2017a; Holoien et al., 2018).
Its X-ray flux increased around $\sim$200 d before decaying $\sim$350 d after
discovery, which is believed to have been caused by either inefficient
circularization that resulted in delayed accretion, or material surround the
TDE becoming optically thin a few months after discovery (Gezari et al.,
2017b; Holoien et al., 2018). The optically-detected TDEs PS18kh (Holoien et
al., 2019a; van Velzen et al., 2019a) and ASASSN-19bt (Holoien et al., 2019b)
have also shown X-ray emission, but only the TDE candidate ASASSN-18ul (Wevers
et al., 2019b) exhibits a X-ray light curve similar to that of ASASSN-18jd.
ASASSN-18ul shows a relatively flat X-ray light curve with small timescale
variations where the X-ray luminosity varies by more than an order of
magnitude ($\sim$1042$-$$10^{43.2}\rm~{}erg~{}s^{-1}$). However, the X-ray
spectrum of ASASSN-18ul is different from that of ASASSN-18jd in that it shows
no obvious power-law component and has negligible hard X-ray flux.

###  7.2 ASASSN-18jd as an AGN

As seen in Figures 15 and 16, the spectra of ASASSN-18jd show emission-line
features common to AGN spectra, like C iii] $\lambda$1909, and perhaps O i
$\lambda$1302, Mg ii $\lambda$2800, [Fe xiv] $\lambda$5303, and [Fe x]
$\lambda$6375\. Additionally, the narrow [O iii] and [N ii] emission line
strengths are such that they may imply some low level of pre-existing AGN
activity. However, the transient nature of the coronal Fe lines and the
relatively high ratio of [Fe x] $\lambda$6375 to [O iii] $\lambda$5007
compared to that seen in NLS1s can also be interpreted originating from the
transient soft X-rays generated by a TDE (Komossa et al., 2009; Wang et al.,
2011; Wang et al., 2012). This is due to the gas in the NLR that produces [O
iii] $\lambda$5007 being further away from the center of the AGN than the
coronal Fe-producing gas, and thus the coronal lines are more sensitive to
rapid flares like TDEs.

The X-ray spectra of AGNs share similar properties with ASASSN-18jd, as AGNs
often show a combination of a $kT\sim 100$ eV blackbody and a flat power law
with $\Gamma\simeq 2$ (Ricci et al., 2017; Auchettl et al., 2018).
Additionally, the variable hardness ratio and the soft X-ray flare near +140 d
are behaviors that are more commonly seen in AGNs than in TDEs (Auchettl et
al., 2018). However, whereas the softer-when-brighter phenomena in AGNs is
attributed to a changing power-law index, ASASSN-18jd shows no such
variability in the power-law index, and instead has a fluctuating blackbody
component.

While AGNs can vary dramatically, the frequency of the large fluctuations
occuring over timescales of under a year is very small. In a survey of SDSS
quasars, MacLeod et al. (2012) found that a change in the $g$ band (or in a
filter of similar wavelength) of $|\Delta m_{g}|>1\rm~{}mag$ in a period of
less than 150 d has an occurrence rate of $P\sim 8\times 10^{-6}$. ASASSN-18jd
peaks in $g$ band with roughly $\Delta m_{g}\simeq 1\rm~{}mag$, but because of
host contamination, the true $\Delta m_{g}$ is likely larger. Furthermore, a
change in NUV magnitudes of $|\Delta m|>3.5\rm~{}mag$, the approximate
difference between the measured brightest Swift UVM2 and the archival GALEX
NUV magnitude, even on the timescale of 1$-$5 years, has an occurrence rate of
$P<2\times 10^{-6}$, the sensitivity limit of the survey in MacLeod et al.
(2012). In our case, the difference between our Swift and archival GALEX
photometry is $\sim$8 years, but it is likely safe to assume that the NUV flux
increased on roughly the same timescale as the optical flux. This makes it
unlikely that ASASSN-18jd can be attributed to “normal” AGN variability.

There are the more dramatic changing-look AGNs (e.g., Shappee et al. 2014;
MacLeod et al. 2016) which can have large photometric variation and also
small, short-timescale fluctuations, similar to ASASSN-18jd. However,
changing-look AGNs have spectra typical of “normal” AGNs, while ASASSN-18jd
does not. Lawrence et al. (2016) found a sample of “hypervariable” AGNs with
$|\Delta m_{g}|>1.5\rm~{}mag$ on timescales of 5$-$10 years, and they
attribute some of these flares to changes in accretion state and some to
microlensing. Graham et al. (2017) found a similar population of hypervariable
AGNs on shorter timescales of 1$-$3 years, which they attribute to
microlensing and stellar activity, like SNe, stellar-mass BH mergers, and even
“slow TDEs”, occuring within the AGNs, rather than intrinsic variability of
the AGNs.

###  7.3 ASASSN-18jd as a rapid turn-on AGN

Recently, there have been discoveries of UV-bright transients at the centers
of galaxies that are either quiescent or contain faint, previously undetected
AGNs. These transients are not consistent with known AGN variability or with
conventional optically-detected TDEs due to their slow fading timescales and
lack of very broad (FWHM$\sim 10^{4}\rm~{}km~{}s^{-1}$) emission features
(e.g., Gezari et al. 2017a; Tadhunter et al. 2017). These “rapid turn-on”
events are similar to changing-look AGNs but occur over much shorter
timescales of a few months and are thought to be brought on by a rapid change
in the accretion state of the AGNs.

Kankare et al. (2017) reported the energetic transient PS1-10adi, which
rapidly brightened by $\sim$2 mag before decaying slowly, smoothly, and
exponentially. Its spectra showed features similar to conventional AGNs, with
relatively narrow Balmer features and strong Fe ii features blueward and
redward of H$\beta$. While PS1-10adi had a strong blue continuum indicative of
a blackbody, the maximum effective temperature was $\sim$1.1$\times
10^{4}\rm~{}K$, much cooler than ASASSN-18jd, though not dramatically cooler
than the effective temperatures of ASASSN-14ae (Brown et al., 2016) and
ASASSN-19bt (Holoien et al., 2019b). Additionally, the light curve of
PS1-10adi showed a smooth decay, whereas ASASSN-18jd shows short-timescale
fluctuations on top of an overall decay.

Frederick et al. (2019) assigned the transient ZTF18aajupnt to a new class of
changing-look LINERs. ZTF18aajupnt changed from a LINER to a NLS1 in $<$100 d
and showed a blue continuum consistent with a blackbody with $T\sim 4.5\times
10^{4}\rm~{}K$. Its spectra, included in Figure 15, showed strong, transient
He ii and coronal Fe lines with relative strengths similar to those seen in
ECLEs. ZTF18aajupnt also showed an increase in soft X-ray flux that began
occurring $\sim$60 d after the initial UV/optical rise. Compared to most TDEs,
it faded relatively slowly (though not as slowly as ASASSN-18jd) and was less
luminous. ZTF18aajupnt and ASASSN-18jd occured around SMBHs of similar mass
($10^{7.6}\rm~{}M_{\odot}$), shared transient He ii and coronal Fe emission,
and had fairly similar UV spectra, but ZTF18aajupnt showed no N iii emission
in its optical spectra or blackbody component in its X-ray spectrum and had
much narrower Balmer emission lines (FWHM $\sim 1000\rm~{}km~{}s^{-1}$).

Two other rapid turn-on events, OGLE17aaj (Gromadzki et al., 2019) and
ASASSN-17cv/AT 2017bgt (Trakhtenbrot et al., 2019), were especially
interesting because of the presence of strong He ii $\lambda$4686, N iii
$\lambda$4640, and N iii $\lambda$4100 emission. OGLE17aaj and ASASSN-17cv had
relatively quick rises to peak brightness over a few months and had UV/optical
SEDs that were well fit by $T\sim 10^{4}\rm~{}K$ blackbodies. They differed
from “normal” TDEs because they faded much more slowly, and the broadest
components of their optical emission lines were relatively narrow, with FWHM
$\sim 2200\rm~{}km~{}s^{-1}$. ASASSN-18jd shares with these events relatively
strong He iii and N iii emission and a slow fading timescale, but the Balmer
features associated with ASASSN-18jd are broader than the features associated
with either OGLE17aaj or ASASSN-17cv.

###  7.4 Summary

We report the discovery and follow-up observations of the luminous nuclear
transient ASASSN-18jd. The most important observed properties of ASASSN-18jd
are:

  * $\bullet$

Archival photometry shows little evidence of strong AGN activity prior to the
transient.

  * $\bullet$

Swift UVOT photometry show the continuum emission to be well-modeled by a
luminous, $T\sim 2.5\times 10^{4}\rm~{}K$ blackbody that fades slowly, with
some short-timescale variability. While, the photometry is also well fit by a
power-law consistent with an accretion disc, the UV spectra from HST show the
SED to be more consistent with a blackbody.

  * $\bullet$

Swift XRT and XMM-Newton data show widely varying, yet overall fading X-ray
emission with $kT\sim 100\rm~{}eV$ blackbody and $\Gamma\sim 1.7$ power-law
components.

  * $\bullet$

Optical spectroscopy show strong, roughly constant Balmer emission, as well as
transient He ii $\lambda$4686, N iii $\lambda$4100, and [Fe x] $\lambda$6375.

On balance, ASASSN-18jd can be understood as either a TDE or as a rapid turn-
on event. As a TDE, ASASSN-18jd would be one of the most luminous and slowest
events known and would challenge the paradigm of TDEs declining rapidly and
smoothly. As a rapid turn-on event, ASASSN-18jd would be representative of a
turn-on to a broad line Seyfert 1. As ASASSN-18jd continues to evolve, future
observations will help us understand the nature of this transient.

## Acknowledgements

We thank the Swift PI, the Observation Duty Scientists,and the science
planners for promptly approving and executing our Swift observations. We thank
the XMM-Newton team for promptly scheduling and executing our TOO
observations. We thank the Las Cumbres Observatory and its staff for its
continuing support of the ASAS-SN project.

ASAS-SN is supported by the Gordon and Betty Moore Foundation through grant
GBMF5490 to the Ohio State University and NSF grant AST-1515927. Development
of ASAS-SN has been supported by NSF grant AST-0908816, the Mt. Cuba
Astronomical Foundation, the Center for Cosmology and AstroParticle Physics at
the Ohio State University, the Chinese Academy of Sciences South America
Center for Astronomy (CASSACA), the Villum Foundation, and George Skestos.

CSK and KZS are supported by NSF grants AST-1515876, AST-1515927, and
AST-1814440. CSK is also supported by a fellowship from the Radcliffe
Institute for Advanced Studies at Harvard University. KAA is supported by the
Danish National Research Foundation (DNRF132). BJS is supported by NSF grant
AST-1908952. SD, SB, and PC acknowledge NSFC 11573003. We acknowledge
Telescope Access Program (TAP) funded by the NAOC, CAS, and the Special Fund
for Astronomy from the Ministry of Finance. MAT acknowledges support from the
DOE CSGF through grant DE-SC0019323. CR acknowledges support from the
CONICYT+PAI Convocatoria Nacional subvencion a instalacion en la academia
convocatoria año 2017 PAI77170080. PJV is supported by the National Science
Foundation Graduate Research Fellowship Program Under Grant No. DGE-1343012.
Support for JLP is provided in part by FONDECYT through the grant 1191038 and
by the Ministry of Economy, Development, and Tourism’s Millennium Science
Initiative through grant IC120009, awarded to The Millennium Institute of
Astrophysics, MAS. MRS is supported by the National Science Foundation
Graduate Research Fellowship Program Under Grant No. 1842400. MG is supported
by the Polish NCN MAESTRO grant 2014/14/A/ST9/00121. The UCSC team is
supported in part by NSF grant AST-1518052, the Gordon & Betty Moore
Foundation, the Heising-Simons Foundation, and by a fellowship from the David
and Lucile Packard Foundation to RJF.

Based on observations made with the NASA/ESA Hubble Space Telescope, obtained
at the Space Telescope Science Institute, which is operated by the Association
of Universities for Research in Astronomy, Inc., under NASA contract
NAS5-26555. These observations are associated with programs GO-14781 and
GO-15312.

Based on data acquired using the Large Binocular Telescope (LBT). The LBT is
an international collaboration among institutions in the United States, Italy
and Germany. LBT Corporation partners are: The University of Arizona on behalf
of the Arizona university system; Istituto Nazionale di Astrofisica, Italy;
LBT Beteiligungsgesellschaft, Germany, representing the Max-Planck Society,
the Astrophysical Institute Potsdam, and Heidelberg University; The Ohio State
University, and The Research Corporation, on behalf of The University of Notre
Dame, University of Minnesota and University of Virginia. This paper uses data
obtained with the MODS spectrographs built with funding from NSF grant
AST-9987045 and the NSF Telescope System Instrumentation Program (TSIP) with
additional funds from the Ohio Board of Regents and the Ohio State University
Office of Research.

Some of the observations reported in this paper were obtained with the
Southern African Large Telescope (SALT). Polish participation in SALT is
funded by grant no. MNiSW DIR/WK/2016/07.

This research has made use of the XRT Data Analysis Software (XRTDAS)
developed under the responsibility of the ASI Science Data Center (ASDC),
Italy. At Penn State the NASA Swift program is supported through contract
NAS5-00136.

Observations made with the NASA Galaxy Evolution Explorer (GALEX) were used in
the analyses presented in this manuscript. Some of the data presented in this
paper were obtained from the Mikulski Archive for Space Telescopes (MAST).
STScI is operated by the Association of Universities for Research in
Astronomy, Inc., under NASA contract NAS5-26555. Support for MAST for non-HST
data is provided by the NASA Office of Space Science via grant NNX13AC07G and
by other grants and contracts.

This publication makes use of data products from the Wide-field Infrared
Survey Explorer, which is a joint project of the University of California, Los
Angeles, and the Jet Propulsion Laboratory/California Institute of Technology,
funded by NASA.

This publication makes use of data obtained from the Weizmann Interactive
Supernova Data Repository (WISeREP, Yaron & Gal-Yam 2012).

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Table 3: Photometry of ASASSN-18jd. All magnitudes are in the AB system and have not been corrected for Galactic extinction or host-subtracted, except for ASAS-SN photometry, which has the observed pre-transient flux subtracted out. Non-detections in ASAS-SN are given as 3$\sigma$ upper limits. Only a portion of the data is shown here. The entire table is published in machine-readable format in the online journal. MJD | Filter | Magnitude | Telescope/Observatory
---|---|---|---
56618.219 | $V$ |  $>$18.11 | ASAS-SN
56801.576 | $V$ |  $>$19.05 | ASAS-SN
56805.575 | $V$ |  $>$18.95 | ASAS-SN
$...$ |  |  |
58589.407 | $i$ | $16.30\pm 0.02$ | Swope
58636.339 | $i$ | $16.32\pm 0.02$ | Swope
58660.258 | $i$ | $16.32\pm 0.02$ | Swope
Table 4: X-ray Observations of ASASSN-18jd. Non-detections by Swift XRT are given as 3$\sigma$ upper limits. Flux is measured from fitting the individual spectra for XMM-Newton, while for Swift XRT, flux measurements assume a power-law spectrum with $\Gamma=1.75$ (see Section 4). Only a portion of the data is shown here. The entire table is published in machine-readable format in the online journal. MJD | Exposure | 0.3$-$10 keV Flux | HR | Telescope
---|---|---|---|---
| [s] | [$10^{-13}$ erg cm-2 s-1] |  |
58266.819 | 28000 | $1.920\pm 0.101$ | $-0.514\pm 0.011$ | XMM-Newton
58280.591 | 21000 | $2.970\pm 0.355$ | $-0.554\pm 0.009$ | XMM-Newton
58249.807 | 2163 | $1.380\pm 0.560$ | $-0.540\pm 0.094$ |  Swift XRT
… |  |  |  |
58646.995 | 477 | $<1.574$ | $-0.048$ |  Swift XRT
58651.436 | 2245 | $<0.614$ | 0.052 |  Swift XRT
58661.400 | 1788 | $<0.775$ | 0.090 |  Swift XRT
Table 5: NUV/Optical Spectroscopic Observations of ASASSN-18jd Date [UT] | MJD | MJD$-t_{0}$ | Telescope | Instruments | Waverange [Å] | Resolution [Å] | Slit Width [′′] | Exposure [s]
---|---|---|---|---|---|---|---|---
2018-05-13 | 58251.4 | +34.0 | du Pont 2.5-m | WFCCD | 3600$-$10000 | 3 | 1.5 | 1x600
2018-05-16 | 58254.1 | +36.7 | SALT 10-m | RSS | 3700$-$9000 | 16 | 1.5 | 1x1600
2018-05-16 | 58254.6 | +37.2 | U of Hawaii 88-in | SNIFS | 3200$-$10000 | 7 | IFU | 1x1200, 1x1600
2018-05-18 | 58256.6 | +39.2 | UH88 | SNIFS |  |  |  | 1x2000
2018-06-11 | 58280.1 | +62.7 | SALT | RSS |  |  |  | 1x1600
2018-06-12 | 58281.6 | +64.2 | UH88 | SNIFS |  |  |  | 1x1200
2018-06-17 | 58286.0 | +68.6 | SALT | RSS |  |  |  | 1x1600
2018-06-26 | 58295.0 | +77.6 | SALT | RSS |  |  |  | 1x1600
2018-07-08 | 58307.5 | +90.1 | Shane 3-m | Kast | 3500$-$10000 | 3 | 2 | 3x500
2018-07-11 | 58310.5 | +93.1 | UH88 | SNIFS |  |  |  | 1x1800
2018-07-15 | 58314.5 | +97.1 | Shane | Kast |  |  |  | 3x600
2018-07-18 | 58318.0 | +100.6 | SALT | RSS |  |  |  | 1x1600
2018-07-23 | 58322.6 | +105.2 | UH88 | SNIFS |  |  |  | 1x1800
2018-08-02 | 58332.6 | +115.2 | SALT | RSS |  |  |  | 1x1600
2018-08-14 | 58344.5 | +127.1 | UH88 | SNIFS |  |  |  | 1x2000
2018-08-18 | 58348.8 | +131.4 | SALT | RSS |  |  |  | 1x1600
2018-09-05 | 58366.3 | +148.9 | UH88 | SNIFS |  |  |  | 1x600, 1x1800
2018-09-25 | 58386.3 | +168.9 | UH88 | SNIFS |  |  |  | 1x1800
2018-10-10 | 58401.3 | +183.9 | UH88 | SNIFS |  |  |  | 1x1800
2018-10-22 | 58413.3 | +195.9 | UH88 | SNIFS |  |  |  | 1x1800
2018-10-31 | 58422.9 | +205.5 | SALT | RSS |  |  |  | 1x1600
2018-11-05 | 58427.1 | +209.7 | LBT 8.4-m | MODS | 3200$-$10000 | 3 | 1.0 | 4x1200
2018-11-06 | 58428.2 | +210.8 | UH88 | SNIFS |  |  |  | 1x1800
2018-12-04 | 58456.2 | +238.8 | UH88 | SNIFS |  |  |  | 1x1800
2018-12-31 | 58483.2 | +265.8 | UH88 | SNIFS |  |  |  | 1x1600
2019-01-03 | 58486.2 | +268.8 | UH88 | SNIFS |  |  |  | 1x1800
2019-01-08 | 58491.2 | +273.8 | UH88 | SNIFS |  |  |  | 1x1200
2019-06-04 | 58638.3 | +420.9 | du Pont | WFCCD |  |  |  | 3x1200
2019-06-27 | 58661.0 | +443.6 | SALT | RSS |  |  |  | 1x1800
2019-06-28 | 58662.4 | +445.0 | LBT | MODS |  |  |  | 3x1200
Table 6: HST-STIS Spectroscopic Observations of ASASSN-18jd Date [UT] | MJD | MJD$-t_{0}$ | Exposure (NUV) [s] | Exposure (FUV) [s]
---|---|---|---|---
2018-05-18 | 58256.8 | +39.4 | 4$\times$169 | 4$\times$169
2018-06-01 | 58270.2 | +52.8 | 6$\times$309 | 6$\times$410
2018-06-28 | 58297.5 | +80.1 | 6$\times$309 | 6$\times$410
2018-07-31 | 58330.9 | +113.5 | 5$\times$375, 5$\times$480 | 5$\times$497, 5$\times$419
2018-08-14 | 58344.5 | +127.1 | 5$\times$375, 6$\times$392, | 3$\times$571, 5$\times$471,
|  |  | 1$\times$370 | 3$\times$646, 1$\times$485
2018-08-31 | 58361.2 | +144.7 | 5$\times$375, 6$\times$392, | 3$\times$571, 5$\times$471,
|  |  | 1$\times$370 | 3$\times$646, 1$\times$485