# Discovery and Rapid Follow-up Observations of the Unusual Type II SN 2018ivc
in NGC 1068

K. A. Bostroem Department of Physics, University of California, 1 Shields Avenue, Davis, CA 95616-5270, USA  S. Valenti Department of Physics, University of California, 1 Shields Avenue, Davis, CA 95616-5270, USA  D. J. Sand Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  J. E. Andrews Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  S. D. Van Dyk Caltech/IPAC, Mailcode 100-22, Pasadena, CA 91125 USA  L. Galbany Departamento de Física Teórica y del Cosmos, Universidad de Granada, E-18071 Granada, Spain  D. Pooley Department of Physics and Astronomy, Trinity University, San Antonio, Texas  Eureka Scientific, Inc., USA  R. C. Amaro Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  N. Smith Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  S. Yang Department of Physics, University of California, 1 Shields Avenue, Davis, CA 95616-5270, USA  Department of Physics and Astronomy Galileo Galilei, University of Padova, Vicolo dell’Osservatorio, 3, I-35122 Padova, Italy  INAF - Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5, I-35122 Padova, Italy  G. C. Anupama Indian Institute of Astrophysics, Koramangala 2nd Block, Bengaluru 560 034, India  I. Arcavi The School of Physics and Astronomy, Tel Aviv University, Tel Aviv 69978, Israel  CIFAR Azrieli Global Scholars program, CIFAR, Toronto, Canada  E. Baron Homer L. Dodge Department of Physics and Astronomy, University of Oklahoma, Norman, OK, USA  P. J. Brown Mitchell Institute for Fundamental Physics and Astronomy, Texas A&M University,College Station, TX 77843, USA  J. Burke  Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA  Las Cumbres Observatory, 6740 Cortona Dr, Suite 102, Goleta, CA 93117-5575, USA  R. Cartier  Cerro Tololo Inter-American Observatory, National Optical Astronomy Observatory, Casilla 603, La Serena, Chile  D. Hiramatsu Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA  Las Cumbres Observatory, 6740 Cortona Dr, Suite 102, Goleta, CA 93117-5575, USA  R. Dastidar Aryabhatta Research Institute of observational sciences, Manora Peak, Nainital 263 001, India  Department of Physics $\&$ Astrophysics, University of Delhi, Delhi-110 007, India  J. M. DerKacy Homer L. Dodge Department of Physics and Astronomy, University of Oklahoma, Norman, OK, USA  Y. Dong  Department of Physics, University of California, 1 Shields Avenue, Davis, CA 95616-5270, USA  E. Egami  Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  S. Ertel  Large Binocular Telescope Observatory, 933 North Cherry Avenue, Tucson, AZ 85721, USA  Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  A. V. Filippenko  Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA  Miller Senior Fellow, Miller Institute for Basic Research in Science, University of California, Berkeley, CA 94720, USA  O. D. Fox Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA  J. Haislip  Department of Physics and Astronomy, University of North Carolina at Chapel Hill, Chapel Hill, NC 27599, USA  G. Hosseinzadeh Center for Astrophysics | Harvard & Smithsonian, 60 Garden Street, Cambridge, MA 02138-1516, USA  D. A. Howell Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA  Las Cumbres Observatory, 6740 Cortona Dr, Suite 102, Goleta, CA 93117-5575, USA  A. Gangopadhyay  Aryabhatta Research Institute of observational sciences, Manora Peak, Nainital 263 001, India  Pt. Ravi Shankar Shukla University, Raipur 492 010, India  S. W. Jha Department of Physics and Astronomy, Rutgers, the State University of New Jersey, 136 Frelinghuysen Road, Piscataway, NJ 08854, USA  V. Kouprianov  Department of Physics and Astronomy, University of North Carolina at Chapel Hill, Chapel Hill, NC 27599, USA  Central (Pulkovo) Observatory of Russian Academy of Sciences, 196140 Pulkovskoye Ave. 65/1, Saint Petersburg, Russia  B. Kumar Indian Institute of Astrophysics, Koramangala 2nd Block, Bengaluru 560 034, India  M. Lundquist  Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  D. Milisavljevic Department of Physics and Astronomy, Purdue University, 525 Northwestern Ave., West Lafayette, IN. 47907, USA  C. McCully Department of Physics, University of California, Santa Barbara, CA 93106-9530, USA  Las Cumbres Observatory, 6740 Cortona Dr, Suite 102, Goleta, CA 93117-5575, USA  P. Milne  Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  K. Misra Aryabhatta Research Institute of observational sciences, Manora Peak, Nainital 263 001, India  Department of Physics, University of California, 1 Shields Avenue, Davis, CA 95616-5270, USA  D. E. Reichart Department of Physics and Astronomy, University of North Carolina at Chapel Hill, Chapel Hill, NC 27599, USA  D. K. Sahu  Indian Institute of Astrophysics, Koramangala 2nd Block, Bengaluru 560 034, India  H. Sai  Physics Department and Tsinghua Center for Astrophysics (THCA), Tsinghua University, Beijing, 100084, China  A. Singh Indian Institute of Astrophysics, Koramangala 2nd Block, Bengaluru 560 034, India  Joint Astronomy Programme, Department of Physics, Indian Institute of Science, Bengaluru 560012, India  P. S. Smith  Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  J. Vinko Konkoly Observatory, Research Centre for Astronomy and Earth Sciences, Hungarian Academy of Sciences, H-1121 Budapest, Hungary  Department of Astronomy, University of Texas at Austin, Austin, TX 78712, USA  X. Wang Physics Department and Tsinghua Center for Astrophysics (THCA), Tsinghua University, Beijing, 100084, China  Y. Wang Key laboratory of Optical Astronomy, National Astronomical Observatories, Chinese Academy of Sciences, Beijing, 100101, China  J. C. Wheeler Department of Astronomy, University of Texas at Austin, Austin, TX 78712, USA  G. G. Williams Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  MMT Observatory, P.O. Box 210065, University of Arizona, Tucson, AZ 85721, USA  S. Wyatt  Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  J. Zhang  Yunnan Observatories, Chinese Academy of Sciences, Kunming 650216, China  Key Laboratory for the Structure and Evolution of Celestial Objects, Chinese Academy of Sciences, Kunming 650216, China  X. Zhang  Physics Department and Tsinghua Center for Astrophysics (THCA), Tsinghua University, Beijing, 100084, China

(Accepted April 7, 2020)

###### Abstract

We present the discovery and high-cadence follow-up observations of SN
2018ivc, an unusual Type II supernova that exploded in NGC 1068 ($D=10.1$
Mpc). The light curve of SN 2018ivc declines piecewise-linearly, changing
slope frequently, with four clear slope changes in the first 30 days of
evolution. This rapidly changing light curve indicates that interaction
between the circumstellar material and ejecta plays a significant role in the
evolution. Circumstellar interaction is further supported by a strong X-ray
detection. The spectra are rapidly evolving and dominated by hydrogen, helium,
and calcium emission lines. We identify a rare high-velocity emission-line
feature blueshifted at $\sim 7800$ $\rm{km~{}s^{-1}}$ (in H$\rm{\alpha}$,
H$\rm{\beta}$, P$\rm{\beta}$, P$\rm{\gamma}$, He i, Ca ii), which is visible
from day 18 until at least day 78 and could be evidence of an asymmetric
progenitor or explosion. From the overall similarity between SN 2018ivc and SN
1996al, the H$\rm{\alpha}$ equivalent width of its parent H ii region, and
constraints from pre-explosion archival Hubble Space Telescope images, we find
that the progenitor of SN 2018ivc could be as massive as 52 $\rm{M_{\odot}}$
but is more likely $\rm{<12}$ $\rm{M_{\odot}}$. SN 2018ivc demonstrates the
importance of the early discovery and rapid follow-up observations of nearby
supernovae to study the physics and progenitors of these cosmic explosions.

supernovae: individual (SN 2018ivc) – supernovae: general

††journal: ApJ††facilities: ARC (DIS), ARIES:ST, ARIES:DFOT, Beijing:2.16m
(BFOSC), CTIO:PROMPT, Gemini:South (Flamingos-2), CXO, HCT (HFOSC), HET
(LRS2), HST (ACS, WFPC2, WFC3), Keck: I (LRIS), II (DEIMOS), Las Cumbres
Observatory (FLOYDS, Sinistro), LBT (MODS), MAST, MMT (Binospec), PO:1.2m,
SALT (RSS), SOAR (Goodman), SO: Bok (SPOL), SO: Kuiper (Mont4K) SO:Super-
LOTIS, UMN:1.52m††software:  astropy (Astropy Collaboration et al., 2013; The
Astropy Collaboration et al., 2018), lcogtsnpipe (Valenti et al., 2016), CIAO
(Fruscione et al., 2006), Sherpa (Freeman et al., 2001), SpectRes (Carnall,
2017)

##  1 Introduction

Single stars with masses greater than $\rm{\sim 8}$ $\rm{M_{\odot}}$ are
thought to explode as core-collapse supernovae (SNe). Type II SNe, those with
hydrogen in their spectra, are empirically classified by the shape of their
light curves. The light curves of Type IIP SNe show a long plateau, $\sim
80$–120 days after explosion, before falling by a few magnitudes over $\sim
20$ days and eventually settling on a decline powered by the radioactive decay
of $\rm{{}^{56}Ni\rightarrow{}^{56}Co\rightarrow{}^{56}Fe}$. On the other
hand, the light curves of historical Type IIL SNe (e.g., SN 1979C, SN 1980K)
decline linearly (in mag day-1) before transitioning to the radioactive decay
phase. While historically these two classes were separated (e.g., Barbon et
al., 1979; Patat et al., 1993, 1994; Arcavi et al., 2012; Faran et al., 2014a,
b), today, with larger samples of light curves, we see that these classes blur
together: there are intermediate objects showing a linear decline like that of
Type IIL SNe, and a clear fall onto the radioactive decay tail like that of
Type IIP SNe (Anderson et al., 2014; Valenti et al., 2015; Galbany et al.,
2016a). Additionally, there appears to be a smooth continuum of slopes during
the hydrogen recombination phase. Throughout this paper we will refer to this
collective class as Type IIP/IIL SNe, and we use the Type IIL-like designation
to indicate objects that are similar to the historical Type IIL SNe and Type
IIP-like to describe SNe that show a clear plateau.

The continuity of Type IIP/IIL SN observational properties makes sense when
one considers the physical mechanism producing the plateau in Type IIP-like
SNe. The progenitors of Type IIP/IIL SNe are massive stars with hydrogen
envelopes. This hydrogen is ionized by the SN shock and as it cools, it
recombines, producing a receding recombination front in the expanding ejecta.
The progenitors of Type IIP-like SNe have large hydrogen envelopes and a
recession rate which matches the expansion rate, producing a constant
luminosity. If the progenitor has experienced more mass loss, then the
hydrogen envelope is less massive. In this case, the canonical picture is that
the photosphere recedes faster, leading to a linearly declining light curve,
the steepness of which depends on the amount of hydrogen in the envelope
(Grassberg et al., 1971; Young & Branch, 1989; Blinnikov & Bartunov, 1993;
Moriya et al., 2016). We note that although this is the standard explanation,
the physical mechanism that produces Type IIL-like SNe is not well understood,
and it is also possible that increased $\rm{}^{56}Ni$ mass or interaction with
circumstellar material (CSM) could produce a linearly decaying light curve.

Over the course of its lifetime, the material lost via winds from the
progenitor star leads to a substantial amount of CSM. The configuration of the
CSM and its density depend on the time of the mass loss, the rate of the mass
loss (which could be steady or episodic), and the symmetry of the mass loss.
Once the star explodes, the radiation, shock wave, and (at later times) ejecta
can interact with the material lost prior to explosion, producing
observational signatures such as narrow emission lines, enhanced luminosity,
and blue colors (Smith, 2014). When narrow or intermediate-width hydrogen
emission lines are observed, the SN is denoted a Type IIn SN (see, e.g.,
Filippenko, 1997, for a review of SNe). Historically, CSM was identified in
SNe either through narrow lines produced by unshocked, photoionized CSM,
intermediate-width lines created by shocked CSM, or by light curves that
deviated from the typical Type IIP/IIL shape and color. Recently, signs of
interaction have been identified beyond these traditional diagnostics. CSM
interaction has been seen in Type IIP/IIL SNe at late phases (Andrews et al.,
2010; Mauerhan et al., 2017), early CSM interaction has been invoked to
explain the early-time light curve of most Type IIP/IIL SNe (Morozova et al.,
2017, 2018), and the presence of CSM material very close to the progenitor has
been used to explain early, narrow, high-ionization features (“flash
spectroscopy” lines) in several Type II SNe (e.g., Quimby et al., 2006; Gal-
Yam et al., 2014; Smith et al., 2015; Khazov et al., 2016; Hosseinzadeh et
al., 2018). Each of these signatures of interaction occurs during a specific
phase of evolution and could be missed without continuous and frequent
observations from explosion through the nebular phase. In order to fully
understand the mass-loss history of their progenitors, it is important to
follow Type IIP/IIL SNe as close to explosion as possible, for as long as
possible, with the highest cadence available.

Here, we present observations of SN 2018ivc, a Type IIL-like SN that exploded
in a complex CSM environment in the well-studied Seyfert 2 galaxy NGC 1068
(M77; see Figure 1), and was discovered by the $D<40$ Mpc SN Survey (DLT40;
Tartaglia et al., 2018). We adopt a distance of $D=10.1$ Mpc ($\mu=30.02$ mag;
Tully et al., 2008) and a recessional velocity of 1037 $\rm{km~{}s^{-1}}$
($z$=0.003793; Huchra et al. 1999) for NGC 1068 in this work. Our near-daily
observations of the SN show the presence of several rapid changes, both in the
spectra and in the light curve, which could be easily missed in undersampled
observations of other Type IIL-like SNe. In Section 2 of this paper, we
discuss recent improvements to the DLT40 pipeline that enabled the very rapid
discovery and follow-up observations of SN 2018ivc. We then describe the
photometric and spectroscopic observations in Section 3, determine the
properties of SN 2018ivc and its host galaxy in Section 4, and discuss the
spectroscopic and photometric evolution of SN 2018ivc in Section 5. In Section
6, we present our search for a progenitor in archival images from the Hubble
Space Telescope (HST) and our identification of a SN in the literature with a
similar evolution. Properties of the progenitor system are considered in
Section 7 and the paper is summarized in Section 8.

##  2 DLT40 Discovery and Rapid Follow-up Campaigns

###  2.1 DLT40 Survey and Recent Improvements

We briefly summarize the relevant aspects of the high-cadence DLT40 SN survey,
the mechanics of which are described in more detail in Yang et al. (2017),
Tartaglia et al. (2018), and Yang et al. (2019). DLT40 is a $\sim 12$ hr
cadence search for SNe, targeting galaxies within $D\lesssim 40$ Mpc, designed
with the goal of discovering $\sim 10$ SNe per year within a day of explosion.
Since December 2017, DLT40 has operated two nearly identical 0.41 m
telescopes, one at Cerro Tololo Inter-American Observatory (CTIO) in Chile and
the other at Meckering Observatory in western Australia. Each telescope
strives to observe the same set of $\sim 400$–600 galaxies each night,
providing the effective $\sim 12$ hr search cadence. The exposure time is 45 s
per field, in a Clear or Open filter, with a typical limiting magnitude of
$r\approx 19$ mag in good sky conditions. As we describe below, the search for
SNe happens in real time, nearly 24 hours a day.

Several recent improvements to the DLT40 pipeline led to the immediate
identification of SN 2018ivc as a true SN, and enabled the rapid follow-up
campaign. First, DLT40 SN candidates are scored in real time using a version
of the pixel-based, random forest machine learning algorithm used by the Pan-
STARRS1 survey (Wright et al., 2015). Once a new, strong SN candidate is
identified by the DLT40 algorithm, an email alert is immediately sent to the
team, and a follow-up observation with a DLT40 telescope can be urgently
requested at the click of a button to verify the new transient. If the
response to the automated email alert is immediate, the entire process from
data taking to automated discovery to confirmation imaging takes approximately
five minutes, and truly brings follow-up response into real time. Once a new
SN has been verified with follow-up imaging, we often trigger a sequence of
intranight DLT40 images, which can elucidate rapid light-curve evolution, as
was the case for SN 2018ivc. Links in the internal DLT40 web pages allow for
nearly automated photometric and spectroscopic triggers of new SNe with the
Las Cumbres Observatory network of telescopes (Brown et al., 2013), and we
always request a target of opportunity (ToO) sequence of multiband Swift
images to probe the early ultraviolet (UV) light-curve evolution. The DLT40
team also announces all verified SN candidates immediately through the
Transient Name Service111https://wis-tns.weizmann.ac.il (TNS) – there are no
proprietary candidates.

###  2.2 SN 2018ivc Discovery and Rapid Follow-up Observations

The initial DLT40 discovery image of SN 2018ivc (which was given an internal
DLT40 designation of DLT18aq) was taken 2018 Nov. 24.04 (UT dates are used
throughout this paper) by the PROMPT5 telescope on CTIO, with a magnitude of
$r=14.65\pm 0.02$. The DLT40 machine-learning algorithm identified a strong SN
candidate in the initial image of the NGC 1068 field, and an email alert was
sent out two minutes later. One of us (R. C. A.) quickly inspected the SN
candidate on the DLT40 internal web site and triggered a second observation,
which was taken 14 min after the discovery image. Once the SN candidate was
confirmed, we took a sequence of PROMPT5 images of the field over the next
$\sim 2.5$ hr, during which the SN brightened from the initial $r=14.65$ mag
to $14.61$ mag. During this time period, we reported the SN candidate to the
community (Valenti et al., 2018) and also triggered the Las Cumbres
Observatory network of robotic telescopes to obtain multiband observations
($UBVgri$) as soon as possible (first images on 2018 Nov. 24.11, 2 hr after
discovery). We obtained our first spectrum with the FLOYDS spectrograph on
2018 Nov. 24.28 and a second spectrum with the BFOSC spectrograph on 2018 Nov.
24.70 (Zhang et al., 2018).

The last nondetection in the NGC 1068 field by the DLT40 team was on 2018 Nov.
19. The field of NGC 1068 had not been observed in the days just prior to the
discovery of SN 2018ivc owing to the lunar angle constraint DLT40 places on
its target fields ($\theta_{\rm Moon}>25^{\circ}$).

Figure 1: A color image of SN 2018ivc and its host galaxy NGC 1068 composed of
multiband observations obtained with Las Cumbres Observatory. The inset shows
a close-up image of the nuclear region of NGC 1068. SN 2018ivc is indicated
with the red guider lines northeast of the nucleus.

##  3 Observations

The photometric and spectroscopic follow-up observations of SN 2018ivc were
obtained and coordinated through Las Cumbres Observatory’s Global Supernova
Project (GSP; PI: D. A. Howell), a key project to collect densely sampled
optical light curves and spectra of nearby and bright SNe (e.g., Szalai et
al., 2019; Andrews et al., 2019).

In addition to the extensive optical dataset, we also obtained early UV and
near-infrared (NIR) photometry and NIR spectroscopy throughout the first $\sim
50$ days, and an early X-ray observation222We also obtained Giant Metrewave
Radio Telescope observations. However, we found issues with the calibration
which rendered the observations unusable..

###  3.1 Photometry

Starting within hours of discovery, high-cadence optical photometric data from
the Las Cumbres Observatory telescope network were acquired in the $UBVgri$
bands for SN 2018ivc; all data were reduced using the lcogtsnpipe software
suite (Valenti et al., 2016) on difference images. Photometric monitoring with
the Las Cumbres Observatory was stopped when we were no longer able to detect
the SN against the bright host background, on 2019 Jan. 21.18 (day 60). These
observations were supplemented with high-cadence photometry from the DLT40
survey, observed with the Open/Clear filters and calibrated to the r band.
These observations were reduced using the DLT40 pipeline (Tartaglia et al.,
2018) and photometry was performed on difference images. DLT40 monitoring was
stopped when the SN disappeared behind the Sun.

We augment these observations with data from the 1.04 m Sampurnanand Telescope
in the BVRI bands (Sagar, 1999), the 1.3 m Devasthal Fast Optical Telescope in
the BVRI bands (Sagar et al., 2012), the 2.01 m Himalayan Chandra Telescope
(HCT) in the BVRI bands (Prabhu & Anupama, 2010), the Mont4K instrument on the
1.55 m Kuiper Telescope in the UBV bands, and the 0.6 m Super-LOTIS telescope
(Williams et al., 2008) in the BVRI bands. Point-spread-function (PSF)
photometry was performed on the original images using the DAOPhot (Stetson,
1987) PyRAF333PyRAF is a product of the Space Telescope Science Institute,
which is operated by AURA, Inc., for NASA package. As the photometry of these
observations was not performed on difference images, imperfect background
subtraction produces more scatter in the light curves. The optical light curve
is presented in Figure 2.

Late-time optical observations were obtained with HST on 2019 July 1 using the
Wide Field Camera 3 (WFC3) UVIS channel (F555W and F814W), as part of the ToO
program GO-15151 (PI: S. Van Dyk). We used Dolphot (Dolphin, 2000, 2016) to
extract the photometry from the FLC frames. A list of all photometric
observations is available in the electronic version of Table 1.

SN 2018ivc was also observed with the Neil Gehrels Swift Observatory (Gehrels
et al., 2004). Observations with the Ultra-Violet Optical Telescope (UVOT;
Roming et al., 2005) were reduced and analyzed using the pipeline for the
Swift Optical Ultraviolet Supernova Archive (SOUSA; Brown et al., 2014), which
includes an arithmetic subtraction of the underlying host-galaxy flux measured
from pre-explosion imaging. The magnitudes use the updated calibration from
Breeveld et al. (2010) and are on the UVOT/Vega system. Optical magnitudes are
not reported because the underlying host galaxy was too bright to correct for
the coincidence loss. Upper limits in uvw2 and uvm2 and detections in uvw1 are
included in Table 1 and plotted in Figure 2.

Table 1: An example of the photometric observations of SN 2018ivc. A complete
machine readable table is included in the online materials.

Observation Date | MJD | Phase | Source | Filter | Magnitude | Magnitude Error
---|---|---|---|---|---|---
UT |  | (day) |  |  | (mag) | (mag)
2018-11-13 02:16:15.16 | 58435.09 | -9.16 | CTIO-Prompt5 | Open | $\rm{<15.26}$ | -
2018-11-14 02:33:08.64 | 58436.11 | -8.14 | CTIO-Prompt5 | Open | $\rm{<15.29}$ | -
2018-11-15 02:24:00.86 | 58437.10 | -7.15 | CTIO-Prompt5 | Open | $\rm{<15.79}$ | -
2018-11-15 05:56:24.00 | 58437.25 | -7.00 | ZTF | g | $\rm{<19.82}$ | -
2018-11-15 07:29:16.80 | 58437.31 | -6.94 | ZTF | g | $\rm{<20.16}$ | -
2018-11-16 03:03:18.72 | 58438.13 | -6.12 | CTIO-Prompt5 | Open | $\rm{<15.66}$ | -
2018-11-17 02:50:24.57 | 58439.12 | -5.13 | CTIO-Prompt5 | Open | $\rm{<15.75}$ | -
2018-11-18 02:02:48.19 | 58440.09 | -4.16 | CTIO-Prompt5 | Open | $\rm{<15.57}$ | -
2018-11-19 02:39:18.43 | 58441.11 | -3.14 | CTIO-Prompt5 | Open | $\rm{<19.36}$ | -
2018-11-20 10:14:52.00 | 58442.43 | -1.82 | ATLAS | o | $\rm{<18.60}$ | -
2018-11-24 00:52:16.32 | 58446.04 | 1.79 | CTIO-Prompt5 | Open | 14.65 | 0.01
2018-11-24 01:14:37.24 | 58446.05 | 1.80 | CTIO-Prompt5 | Open | 14.66 | 0.01
2018-11-24 01:06:45.50 | 58446.05 | 1.80 | CTIO-Prompt5 | Open | 14.68 | 0.01
2018-11-24 01:33:23.90 | 58446.06 | 1.81 | CTIO-Prompt5 | Open | 14.63 | 0.01
2018-11-24 02:14:47.04 | 58446.09 | 1.84 | LCO LSC 1m | U | 14.35 | 0.02

  * The UVW1, UVW2, UVM2, U, B, and V filters are given in the Vega magnitude system; the g, r, and i filters are given in the AB magnitude system.

Figure 2: The light curve of SN 2018ivc at UV (square symbols) and optical
(circles and pentagons) wavelengths. Difference-image photometry is indicated
by a black outline. Upper limits are denoted with arrows. All phases are
calculated with respect to the inferred explosion epoch (2018 Nov. 22.25; see
Section 4) and the period prior to this time shaded in gray. The piecewise
linear fit of the DLT40 r-band light curve is shown in pink. Individual
filters have been offset by a constant (denoted in the legend) for ease of
viewing. The black dashed line shows the V-band slope expected if the light
curve is fully powered by the radioactive decay of 56Co.

###  3.2 Spectroscopy

Spectroscopic observations from a variety of telescopes and instruments were
obtained almost daily in the optical, starting within 5 hr of discovery and
continuing through day 35. Less frequent monitoring with larger telescopes was
performed until SN 2018ivc disappeared behind the Sun $\sim 80$ days post-
discovery. A single nebular spectrum was obtained with Keck/DEIMOS on day 279
(see Figure 3).

Figure 3: The Keck/DEIMOS nebular spectrum, observed on day 279 (black; pink
is smoothed by 25 pixels) with a zoom-in on the H$\rm{\alpha}$ region in the
inset. A single Gaussian fit is shown in the inset in blue. Although the
spectrum has a low signal-to-noise ratio, there is no evidence of multiple
components in the H$\rm{\alpha}$ feature.

Similarly, NIR observations began on day 3 and continued almost weekly through
day 53. A selection of spectra is shown in Figure 4 and all spectroscopic
observations are listed in Table LABEL:tab:spec. All of the spectra presented
in this work were obtained at the parallactic angle to minimize atmospheric
refraction (Filippenko, 1982). These spectra will be made available on
WISeREP444http://wiserep.weizmann.ac.il (Yaron & Gal-Yam, 2012).

Optical spectra were reduced using standard techniques, including bias
subtraction, flat fielding, and cosmic-ray rejection. SN 2018ivc is embedded
in NGC 1068; given this complex background, local sky subtraction was very
important for the spectral extraction. Despite the care taken, some narrow
emission lines are still visible in the final reduced spectra. In these cases,
after visual inspection of two-dimensional spectra and our highest resolution
data (see Section 5.2 for a detailed discussion), we believe these originate
from the host galaxy.

Flux calibration was performed using standard-star observations. For the NIR
spectra, the data were reduced in a similar manner as by Hsiao et al. (2019),
using the standard ABBA technique; observations were taken of nearby A0V stars
adjacent to the science exposures to facilitate telluric corrections and flux
calibration (e.g., Vacca et al., 2003).

Figure 4: Left: the spectroscopic evolution of SN 2018ivc at optical
wavelengths. SN 2018ivc shows strong emission lines and very little
absorption. Imperfect background subtraction may create artificially blue
continua and/or strong narrow lines in the broad H$\rm{\alpha}$ emission. The
phase of each spectrum is given on the right. High resolution observations are
resampled with SpectRes (Carnall, 2017) to 2Å resolution to improve the
signal-to-noise ratio. Right: the NIR spectroscopic evolution with the phase
of the spectrum given on the right. The strong telluric absorption between
13450 Å and 14000 Å is masked in all NIR spectra. At the top of each panel
prominent lines are identified, and in the left panel a prominent telluric
feature is marked with a cross.

###  3.3 X-ray Observations

The Chandra X-ray Observatory observed SN 2018ivc on 2018 Dec. 05.7 for 10.0
ks (ObsID 20306) with the telescope aimpoint on the Advanced CCD Imaging
Spectrometer (ACIS) S3 chip as part of a program to follow up possible X-ray
detections from other facilities (PI: D. Pooley). The host galaxy, NGC 1068,
has been observed many times previously with Chandra, often but not always
with the High Energy Transmission Gratings (HETG) in place. For a
straightforward comparison to our observation, we selected the longest ACIS-S3
observation with no grating in the Chandra data archive: ObsID 344 (PI: A.
Wilson) began on 2000 Feb 21.7 and had an exposure time of 47.4 ks. Data
reduction was performed with the chandra_repro script, part of the Chandra
Interactive Analysis of Observations (CIAO) software (Fruscione et al., 2006).
We used CIAO version 4.11 and calibration database (CALDB) version 4.8.3.

The source is clearly detected (Figure 5) with 207 total counts recorded in a
1$\farcs$5 radius source extraction region in the 0.5–8 keV band. The
background contribution to this X-ray flux is non-negligible but difficult to
estimate given the non-spatially-uniform X-ray emission immediately
surrounding the location of SN 2018ivc in its host galaxy. One estimate of the
background comes from an annular background region of inner radius 2″ and
outer radius 4″ centered on the SN. Based on the 80 counts in this background
region, there are $192\pm 14$ net counts from SN 2018ivc in the 0.5–8 keV
band.

Figure 5: Pre-SN (left) and post-SN (right) images of NGC 1068 taken with
Chandra/ACIS-S3. Each image is 1′ on a side. The white circle on the right is
1$\farcs$5 in radius centered on SN 2018ivc and is used to extract the source
spectrum. Note the non-uniform extended emission surrounding the SN. The cyan
circle on the left (identical to the source region) and the cyan annulus on
the right were each used to extract background spectra. See text for details.

To assess the spectral properties of the SN and determine its X-ray flux, we
extract and simultaneously fit source and background spectra in the 0.5–8 keV
band using Sherpa (Freeman et al., 2001) with the modified Cash (1979)
statistic cstat and the simplex optimization method. Our model components in
all cases are absorbed hot plasmas (APEC model). We use the Tuebingen-Boulder
Interstellar Medium absorption model (Wilms et al., 2000) with a minimum
column density equal to the Galactic value of
$n_{\mathrm{H}}={1.55}\\!\times\\!10^{19}~{}\mathrm{cm}^{-2}$. We separately
use two options for the background spectrum: the first is the annulus
mentioned above in ObsID 20306 and the second is a 1$\farcs$5 radius circle
(identical to the source extraction region) in ObsID 344.

The first choice of background spectrum (annular extraction region from ObsID
20306) is well fit by two hot plasmas (temperatures of $kT_{1}=0.01$ keV and
$kT_{2}=0.85$ keV) absorbed by a column of density ${3.2}\\!\times\\!10^{20}$
$\mathrm{cm}^{-2}$. The second choice of background spectrum (circular
extraction region from ObsID 344) is also well fit by two hot plasmas
(temperatures of $kT_{1}=0.21$ keV and $kT_{2}=0.95$ keV) absorbed by a column
of density ${1.6}\\!\times\\!10^{19}$ $\mathrm{cm}^{-2}$. Using each of these
background spectra separately, we fit the SN spectrum with an absorbed plasma.

With the first choice of background, we obtain a good fit (reduced cstat of
0.63) with best-fit parameters for the SN of
$n_{\mathrm{H}}={(4.4^{{}_{+0.7}}_{{}^{-0.9}})}\\!\times\\!10^{22}~{}\mathrm{cm}^{-2}$
and $kT=17^{{}_{+63}}_{{}^{-7}}$ keV with the upper limit representing the
model maximum. All uncertainties are 68% confidence intervals. The intrinsic
X-ray flux of the SN in this model is ${(8.2\pm 0.8)}\\!\times\\!10^{-13}$
$\mathrm{erg~{}cm}^{-2}~{}\mathrm{s}^{-1}$. For a distance of 10.1 Mpc, this
corresponds to a luminosity of $L_{\mathrm{x}}={(1.0\pm
0.1)}\\!\times\\!10^{40}~{}\mathrm{erg~{}s}^{-1}$.

With the second choice of background, we also obtain a good fit (reduced cstat
of 0.59) with best-fit parameters for the SN of
$n_{\mathrm{H}}={(3.6^{{}_{+0.8}}_{{}^{-0.6}})}\\!\times\\!10^{22}~{}\mathrm{cm}^{-2}$
and $kT=43^{{}_{+37}}_{{}^{-31}}$ keV with the upper limit representing the
model maximum. The intrinsic X-ray flux of the SN in this model is ${(7.5\pm
0.7)}\\!\times\\!10^{-13}$ $\mathrm{erg~{}cm}^{-2}~{}\mathrm{s}^{-1}$. For a
distance of 10.1 Mpc, this corresponds to a luminosity of
$L_{\mathrm{x}}={(9.2\pm 0.9)}\\!\times\\!10^{39}~{}\mathrm{erg~{}s}^{-1}$.

Although the different choices for background extraction give slightly
different results, they are consistent with each other within the
uncertainties. In each case, the reported fluxes are integrated from the
unabsorbed models. Uncertainties on those fluxes are calculated as the
68%-confidence bounds of the integrated, unabsorbed fluxes from Monte Carlo
realizations (1000 samples) of the best-fit models, taking into account the
uncertainties in the best-fit parameters (using the sample_flux command in
Sherpa)

##  4 Properties of the Supernova and its Host Galaxy

###  4.1 Supernova Parameters

The DLT40 survey identified SN 2018ivc on 2018 Nov. 24.07, 4.96 days after the
last observation of the field on 2018 Nov. 19.11. After the DLT40 team’s
prompt reporting of the SN to TNS, the Asteroid Terrestrial-impact Last Alert
System (ATLAS) identified a more recent nondetection in their data on 2018
Nov. 20.42 with a limiting magnitude of 18.6 mag in the orange-ATLAS filter.
As an explosion epoch we adopt 2018 Nov. $22.25\pm 1.8$, the midpoint of the
last nondetection (by ATLAS) and the first detection (by DLT40).

We adopt Milky Way extinction values of $E(B-V)=0.0289\pm 0.0004$ mag from
Schlafly & Finkbeiner (2011). Unfortunately, SN 2018ivc exploded in a region
of high extinction within NGC 1068, which prevented us from deriving the
extinction in the host galaxy from the Na i D lines (Poznanski et al., 2012).
Instead, we estimate an extinction of $E(B-V)=0.5\pm 0.15$ mag by matching the
color evolution of SN 2018ivc over the first 20 days with that of other Type
II SNe (SN 1980K: Barbon et al. 1982; Buta 1982; Tsvetkov 1983; SN 1998S:
Fassia et al. 2000; Liu et al. 2000; Li et al. 2011; Pozzo et al. 2004; SN
1996al: Benetti et al. 2016; SN 2012A: Tomasella et al. 2013a, SN 2013by:
Valenti et al. 2015, SN 2013ej: Valenti et al. 2014). While the light curves
of these SNe have varying slopes, the colors are consistent during the first
20 days, which is why we choose this period to constrain the extinction. A
conservative error of 0.15 mag is adopted to take into account the large
uncertainty of the method. This is consistent with the extinction derived in
the next section from spectroscopy of the parent H ii region and we adopt it
as the host-galaxy extinction for our analysis. Throughout this paper, unless
otherwise noted, we use the extinction law of Cardelli et al. (1989) with
$R_{V}=3.1$.

In Type IIP/IIL SNe, the amount of $\rm{{}^{56}Ni}$ synthesized in the
explosion can be calculated from the luminosity of the SN after the fall from
plateau ($\sim 80$–100 days post explosion), when the light curve is powered
by the radioactive decay of
$\rm{{}^{56}Ni\rightarrow{}^{56}Co\rightarrow{}^{56}Fe}$. We measure the
pseudo-bolometric luminosity of SN 2018ivc from the HST observations on day
220.9 by first transforming the F555W filter and the F814W filter to the
Landolt V and I filters, respectively, using the relations of Harris (2018).
This first step is necessary as SN 1987A was not observed in the WFC3 F555W
and F814W filters. Following Valenti et al. (2008), we find the pseudo-
bolometric luminosity by integrating the apparent magnitude at the effective
wavelength of each filter, using Simpson’s rule. We calculate the pseudo-
bolometric luminosity for SN 1987A from the V and I filters in the same way.
Then, following Spiro et al. (2014), we calculate the nickel mass as
$M({}^{56}{\rm Ni})=0.075\,M_{\odot}\times{L_{\rm 18ivc}(t)}/{L_{\rm
87A}(t)}$. We find $M(\rm{{}^{56}Ni)=0.0056^{+0.0036}_{-0.0022}}\,M_{\odot}$.
Uncertainties in the nickel mass were calculated using a Monte Carlo
simulation taking into account normal uncertainties in the explosion epoch,
distance modulus, Galactic extinction, host galaxy extinction, and apparent
magnitude of SN 2018ivc and SN 1987A. We caution that to calculate this value
we assumed that there was complete $\gamma$-ray trapping, that we could
convert HST filter magnitudes to Landolt filter magnitudes using relations
derived from stellar spectral energy distributions (SEDs), and that SN 1987A
and SN 2018ivc have the same SED. The uncertainties associated with each of
these assumptions are not included in the reported uncertainties.

###  4.2 Host Properties

We searched in the ESO Science
Portal555http://archive.eso.org/scienceportal/home for MUSE integral field
unit (IFU) observations of SN 2018ivc’s host galaxy, NGC 1068. It was observed
on 2014 Dec. 14 under program 094.B-0298(A), in four pointings which we
combined for a total exposure time of 1180 s. All observations were reduced
with the standard MUSE pipeline (Weilbacher et al., 2014) using default
parameters through reflex (Freudling et al., 2013).

For each spaxel, we performed a similar analysis to that of Galbany et al.
(2014, 2016b, 2016c). Briefly, using a modified version of STARLIGHT (Cid
Fernandes et al. 2005; López Fernández et al. 2016, priv. comm.), we model the
stellar component of the continuum by estimating the fractional contribution
of simple stellar populations (SSP) from the Bruzual (2007) base adding dust
attenuation effects as a foreground screen with a Fitzpatrick (1999) reddening
law and $R_{V}=3.1$. Our basis set is composed of 66 SSPs with 17 ages,
ranging from 1 Myr to 18 Gyr, and four different metallicities (0.2, 0.4, 1.0,
and 2.5 $Z_{\odot}$).

By subtracting the best SSP fit from each observed spectrum, we obtained a
pure gas emission spectrum for each spaxel. From the pure gas spectrum, we
estimated the flux of the most prominent emission lines after correcting for
dust content with a correction derived from the Balmer decrement (assuming
case B recombination; Osterbrock & Ferland 2006, the same extinction law, and
$R_{V}=3.1$). From the pure gas models for each spaxel, we create an
extinction-corrected H$\rm{\alpha}$ map.

We use our extinction-corrected H$\rm{\alpha}$ map to identify the H ii region
containing SN 2018ivc (its “parent” region) and derive environmental and
progenitor properties. Following Galbany et al. (2016d, 2018a), the
H$\rm{\alpha}$ maps are used to select star-forming H ii regions across the
galaxy with a modified version of
HIIexplorer666http://www.caha.es/sanchez/HII_explorer/ (Sánchez et al., 2012),
a package that detects clumps of higher intensity in a map by aggregating
adjacent pixels. This procedure selected 1801 H ii clumps with an average
radius of 140 pc. Once the H ii regions were identified, the same analysis
described above was performed on the extracted spectra. The observed spectrum
and the best-fit STARLIGHT spectrum of the H ii region containing SN 2018ivc
can be seen in the bottom-left panel of Figure 6. From this analysis and using
the Cardelli et al. (1989) extinction law, we find $E(B-V)=0.37\pm 0.04$ mag
for the parent region of SN 2018ivc, consistent with the value derived in the
previous section.

We use oxygen abundance as a proxy for metallicity as oxygen is produced at
the beginning of the enrichment process by massive stars and, by mass,
comprises 50% of the heavy elements in the Universe (López-Sánchez et al.,
2012). Using the pure gas emission spectrum of the H ii region, we determine
metallicity from the O3N2 empirical calibrator (Pettini & Pagel, 2004). This
same spectrum is used to find the equivalent width of H$\rm{\alpha}$
(EWH$\rm{\alpha}$). The upper-left panel of Figure 6 shows the
EWH$\rm{\alpha}$ values for each H ii region in NGC 1068. We find the
metallicity to be $\rm{12+log_{10}(O/H)=8.6\pm 0.26}$ and the
EWH$\rm{\alpha}$$=38.58\pm 0.26$ Å. These host properties will be put in the
context of other SNe in Section 7.

Figure 6:  Upper left: the map of the H ii regions in NGC 1068 identified by
HIIexplorer; the parent H ii region of SN 2018ivc is marked with a red point.
The map is colored by H$\rm{\alpha}$ equivalent width (EWH$\rm{\alpha}$),
which is a proxy for regions with young stellar populations. A close-up view
of a 10 ″region around the SN is shown in the inset. Lower left: the MUSE
spectrum of the parent H ii region of SN 2018ivc (black) and the STARLIGHT fit
to the spectrum in pink. Upper right: the cumulative fraction of SNe from the
PISCO sample (Galbany et al., 2018b) as a function of metallicity for Type
IIP/IIL SNe (black) and Type IIn SNe (pink). The metallicity of the parent H
ii region of SN 2018ivc is marked with a blue vertical line and the
uncertainty in this value is denoted with the light blue shaded region. Solar
metallicity (8.69; Asplund et al. 2009) is plotted as a dotted yellow line.
Lower right: the cumulative fraction of SNe as a function of EWH$\rm{\alpha}$
of their parent H ii region. The fraction of Type IIP/IIL SNe is shown in
black, the fraction of Type IIn SNe in pink, and the EWH$\rm{\alpha}$ of SN
2018ivc is marked with the blue vertical line.

##  5 Supernova Evolution

The detailed evolution of SN 2018ivc was caught in the almost daily
photometric and spectroscopic observations. These revealed a rapidly evolving
SN with a steeply declining light curve and spectra with broad emission.

###  5.1 Light-Curve Evolution

The light curve of SN 2018ivc rose rapidly, peaking around day 3, only 5 days
after the last nondetection. It then showed a rapid, linear decline typical of
Type IIL-like SNe. Figure 7 displays a comparison of SN 2018ivc and other
well-studied SNe with a variety of slopes. While the global trend is linear
with a similar decline rate to SNe 1979C and 1980K, our high-cadence
observations quickly revealed that there are, in the first 50 days, several
changes in the slope rarely seen in other Type IIL-like SNe. SN 2018ivc is
also relatively faint compared to other SNe with similar slopes.

Figure 7: The light curve of SN 2018ivc (black stars; DLT40 r band) compared
to a sample of well-studied SNe. SNe with a variety of slopes are represented
with an emphasis on Type IIL-like SNe. The best-observed filter is shown for
each SN. SN 1979C (pink circles, V band; Balinskaia et al., 1980; de
Vaucouleurs et al., 1981; Barbon et al., 1982) and SN 1980K (blue squares, V
band; Barbon et al., 1982; Buta, 1982; Tsvetkov, 1983) represent the
historical Type IIL-like class. SN 2017eaw (sea foam green points, V band;
Szalai et al., 2019) represents the Type IIP-like SNe. SN 2013ej (light blue
triangles, V band; Valenti et al., 2014) and SN 2014G (maroon arrows, V band;
Terreran et al., 2016) are transitional Type IIL-like objects, with a clear
fall from the plateau onto the radioactive decay tail. We plot two SNe that
will be discussed later in the text in comparison to SN 2018ivc: SN 2010jp
(green arrows, R band; Smith et al., 2012) and SN 1996al (yellow pentagons, R
band; Benetti et al., 2016).

The light curve of SN 2018ivc changes slope approximately every 10 days over
the first 30 days. We fit a continuous piecewise linear function to the r-band
observations, leaving the slopes, initial intercept, and breakpoints as free
parameters (8 free parameters). This fit is shown in Figure 2. The initial
decline ending at day $\rm{7.57\pm 0.41}$ has a slope of $\rm{0.1068\pm
0.0094}$ mag $\rm{day^{-1}}$. This is followed by a plateau until day
$\rm{18.07\pm 0.45}$ with a slope of $\rm{0.0056\pm 0.0031}$ mag
$\rm{day^{-1}}$. The light curve begins to decline steeply on day
$\rm{18.07}$, with a slope of $\rm{0.0811\pm 0.0077}$ mag $\rm{day^{-1}}$.
After day $\rm{27.55\pm 1.37}$, the slope changes again to $\rm{0.0355\pm
0.0014}$ mag $\rm{day^{-1}}$. At no point is the slope of SN 2018ivc
consistent with the slope expected for cobalt decay (0.008 mag $\rm{day^{-1}}$
in $R$). Searching among other Type IIP/IIL SNe, we were only able to find one
light curve in the literature that resembled the evolution (although $\sim 1$
mag brighter): that of SN 1996al (see Section 6.1).

We searched the DLT40 difference images in the months leading up to SN 2018ivc
for a pre-explosion outburst similar to SN 2009ip ($M_{r}\approx-14.5$ mag;
Foley et al. 2011; Mauerhan et al. 2013; Pastorello et al. 2013; Margutti et
al. 2014). However, these observations do not reveal any hint of a pre-
outburst eruption, with $\sim 90$ images of the field take in the $\sim 150$
days prior to the SN explosion, with a typical limiting magnitude of $r\approx
19.3$ ($M_{r}\approx-12$ mag).

###  5.2 Spectroscopic Evolution

The spectroscopic evolution of SN 2018ivc shows strong H and He i emission and
very shallow (if any) absorption. The full evolution can be seen in Figure 4,
where prominent spectroscopic features are labeled.

The optical and NIR spectra obtained at 2.04 and 2.95 days (respectively) show
a featureless blue continuum. By day 5, broad H$\rm{\alpha}$ and He i $\lambda
5876$ emission begin to develop with similar profiles. While the presence of
H$\rm{\beta}$ in absorption cannot be completely ruled out (because of
possible contamination from other lines), an absorption component is clearly
missing from both H$\rm{\alpha}$ and He i $\lambda 5876$. This is typical of
Type IIn SNe and some Type IIL-like SNe (Gutiérrez et al., 2014). Although the
physical mechanism for the lack of absorption is still unclear, possible
explanations include a low-density (possibly low-mass) envelope and scattering
off of CSM (Schlegel, 1996).

The Ca II $\lambda\lambda 8498$, 8542, 8662 emission triplet begins to develop
at day 12 and strengthens through day 78. On day 30 the He i $\lambda 7065$
line begins to show in emission. By day 75, the He i $\lambda 6678$ line is
visible in the red wing of the H$\rm{\alpha}$ emission; this feature may be
blended in earlier spectra. The shape of the emission around H$\rm{\alpha}$ in
the two spectra from days 75 and 78 is very boxy, which we will discuss more
in later sections. We obtained a low signal-to-noise ratio (S/N) spectrum at
day 279 which shows a broad H$\rm{\alpha}$ profile, similar to that seen on
day 78.

The NIR spectra are dominated by hydrogen and He i emission features. The blue
continuum at day 3 gives way to Paschen emission features by day 10. The He i
$\lambda 10830$ line is blended in the P$\rm{\gamma}$ feature. In both the
optical and the NIR, the widths of the emission lines decrease with time as
the speed of the ejecta decreases.

The spectra of SN 2018ivc show potential narrow emission lines on top of the
broad H$\rm{\alpha}$ emission. This region is heavily contaminated by the
host-galaxy emission which is challenging to separate from the H$\rm{\alpha}$
SN emission in all but the highest resolution observations. We searched for a
SN component in our highest resolution spectrum: the MMT/Binospec observation
from 2018 Dec. 14. In this spectrum we fit a Gaussian profile to night-sky
lines, galaxy emission lines, and H$\rm{\alpha}$. We find the full width at
half-maximum intensity (FWHM) of the narrow H$\rm{\alpha}$ emission ($\sim
130$ km $\rm{s^{-1}}$) to be consistent with that of the other galaxy emission
lines ([N II] $\lambda\lambda 6548$, 6583; [S II] $\lambda\lambda 6716$,
6731).

We measure the velocity evolution using the half width at half-maximum
intensity (HWHM) of Gaussian fits to the H$\rm{\alpha}$, He i $\lambda 5876$,
and P$\rm{\beta}$ emission lines and the minimum of the Fe ii and
H$\rm{\beta}$ absorption features. The velocities are shown in Figure 8, with
velocities measured from emission features in the left panel and velocities
measured from absorption features in the right panel. We note that while there
is asymmetry, substructure, and host contamination in individual features, the
trends are consistent across features, giving us confidence in the overall
velocity evolution. However, we caution that there may be significant scatter
in individual measurements. The H$\rm{\alpha}$ velocity can be compared to the
average of 112 Type IIP/IIL SNe, calculated by Gutiérrez et al. (2017). We
find that the velocities are similar to those found in Type IIP/IIL SNe and
higher than the velocities of SN 1996al. However, the evolution is more rapid,
with the initial velocity higher than average and the final velocities lower
than average. The Fe ii lines are also broadly consistent with the average
Type IIP/IIL evolution although there is considerable scatter. We
simultaneously fit Fe ii $\lambda\lambda 4924$, 5018, but we were unable to
simultaneously fit Fe ii $\lambda\lambda 4924$, 5018, 5169 and therefore fit
Fe ii $\lambda 5169$ separately. Velocities from both fits are shown in Figure
8.

Figure 8: Left: the expansion velocity evolution of SN 2018ivc found using the
HWHM of emission lines H$\rm{\alpha}$ (pink squares), He i (blue pentagons),
and P$\rm{\beta}$ (yellow diamonds). The H$\rm{\alpha}$ velocity of SN 1996al
is shown as green circles. The velocity of SN 2018ivc is similar to that of
other Type IIP/IIL SNe, although it evolves more rapidly than most, including
SN 1996al. Right: the expansion velocity evolution of SN 2018ivc found using
the minimum wavelength of the absorption features H$\rm{\beta}$ (pink squares)
and Fe ii (blue and yellow circles). The average H$\rm{\alpha}$ velocity
(left) and Fe ii (right) of the Gutiérrez et al. (2017) sample of Type IIP/IIL
SNe is shown in black with the standard deviation in gray.

On day 18 an emission feature begins to develop in the blue wing of the
emission-line profile of H$\rm{\alpha}$, H$\rm{\beta}$, He i $\lambda 5876$,
Ca ii triplet, He i $\lambda 10830$, and P$\rm{\beta}$ (see the first five
panels of Figure 9). This feature gains strength until day $\approx 35$, after
which it fades until it is barely visible on day 78. We identify the emission
line that appears blueward of the rest wavelength of these features as a high-
velocity (HV) component owing to its presence at the same velocity in each
feature (see right panel of Figure 9). From the blueshifted peak of the
emission we find the velocity of the emitting material to be $\sim 9000$
$\rm{km~{}s^{-1}}$ at day 18, slowing down to $\sim 7000$ $\rm{km~{}s^{-1}}$
by day 78.

Figure 9: The evolution of the high-velocity (HV) features in hydrogen
(H$\rm{\alpha}$, H$\rm{\beta}$, P$\rm{\beta}$), helium ($\lambda 5876$,
$\lambda 10830$), and the Ca ii triplet from day 18 through day 78 in the
optical and day 21 through day 53 in the NIR. The feature is most prominent in
H$\rm{\alpha}$, P$\rm{\beta}$, and He i $\lambda 10830$. The panels are
labeled in the upper-left corner with the feature they present. The yellow
shaded region marks $-10,000$ to $-6000$ $\rm{km~{}s^{-1}}$. The phases of the
optical spectra are marked to the left of the first panel, and the NIR phases
are marked to the left of the He i $\lambda 10830$ panel. The Ca ii panel is
centered on the $\lambda 8662$ component of the triplet (marked in yellow at
both rest wavelengths and shaded in yellow at HV). The two blue components are
marked at rest wavelength with vertical cyan lines and at high velocity with a
shaded cyan region. The last panel shows the optical features on day 30 and
the NIR features on day 34 for all emission lines in which the HV feature is
seen. The HV feature is marked with a dashed line at $-8000$
$\rm{km~{}s^{-1}}$ and each emission feature is labeled to the right of the
panel. The bluest two Ca ii triplet features that are offset in velocity space
are marked with a solid line gray at the velocity corresponding to their rest
wavelengths and dashed gray line for the HV features.

##  6 Analysis

###  6.1 Comparison to 1996al

Figure 10 (left panel) shows the DLT40 r-band light curve of SN 2018ivc in
black compared with the R-band light curve of SN 1996al in pink, shifted by
1.0 mag. Since there are no stringent detection limits to constrain the
explosion epoch for SN 1996al, we use an explosion epoch of 1996 July 19,
which is 18 days later than the reference epoch suggested by Benetti et al.
(2016). They compare the light curve to that of other Type IIL-like SNe and
estimate 1996 July 1 as the V-band maximum, which they adopt as the reference
epoch. However, given the similarity between SN 1996al and SN 2018ivc, using
the latter to constrain the explosion of SN 1996al may be more appropriate.
While SN 2018ivc is very well sampled, the sparse sampling of SN 1996al,
especially after the initial decline, makes a detailed comparison challenging.
Nevertheless, the light curve of SN 1996al does show a relatively steep
decline, at a similar rate as SN 2018ivc that, like SN 2018ivc, is interrupted
by a short plateau.

Figure 10: Comparison of the photometric and spectroscopic evolution of SN
2018ivc to the similar event SN 1996al. Left: The r-band light curve of SN
2018ivc (black) compared with the R-band light curve of SN 1996al (pink). The
light curve of SN 1996al has been shifted down by 1.0 mag to align with the SN
2018ivc light curve and the explosion epoch is set to 1996 July 19. The
piecewise linear fit of the SN 2018ivc light curve is shown in blue. The
dashed vertical lines correspond to the phase of the spectra plotted on the
right. Right: The spectroscopic evolution of SN 2018ivc (colored spectra)
compared with that of SN 1996al (gray spectra). The phases of SN 2018ivc are
given on the right side of the figure in color and the phases of SN 1996al are
given on the left side in gray. The spectra of SN 1996al have been offset to
align with those of SN 2018ivc. Although SN 1996al does not show the high-
velocity features and SN 2018ivc does not have narrow lines (the host
contamination visible in Figure 4 has been masked out of the spectra of SN
2018ivc), SN 2018ivc and SN 1996al are spectroscopically similar with strong
He i and hydrogen features in emission.

A spectroscopic comparison of SN 1996al (gray) and SN 2018ivc (color) is shown
in the right panel of Figure 10. As described above, the phases are with
respect to a reference time of 1996 July 19 for SN 1996al. There is broad
agreement in the spectral evolution of both SNe, with H and He i emission
lines dominating the spectra and little to no absorption. The width of the
emission features of SN 1996al are smaller than those of SN 2018ivc, implying
a higher ejecta velocity for SN 2018ivc. This is quantified in the left panel
of Figure 8, which shows that the velocity of H$\rm{\alpha}$ is slower in SN
1996al than in SN 2018ivc and that the H$\rm{\alpha}$ velocity declines more
rapidly than that of SN 1996al.

Given the overall similarity between SN 2018ivc and SN 1996al, we will here
summarize the main results of Benetti et al. (2016) for SN 1996al and compare
these two objects more closely. SN 1996al was identified as a transition SN
between a Type IIL-like and Type IIn with a linearly declining light curve and
a week-long plateau starting around day 15. The spectra show broad hydrogen
and He i emission with narrow P Cygni lines superimposed. From ground-based
pre-explosion H$\rm{\alpha}$ images, a 25 $\rm{M_{\odot}}$ luminous blue
variable (LBV) progenitor was identified. Light-curve modeling of the 15-year
evolution showed that the light curve is dominated by ejecta-CSM interaction.
The linear decline of the light curve, the velocity evolution, and the
evolution of the H$\rm{\alpha}$ flux a year after explosion indicated a low
ejecta mass and that the ejecta were predominantly helium. Benetti et al.
(2016) interpret the similarity in shapes of the hydrogen and helium emission-
line profiles as evidence that both these lines originate in the interaction
region during at least the first 50 days of evolution. The low velocities and
multicomponent emission lines point to a dense, asymmetric CSM, while the
narrow P Cygni lines indicate a patchy, symmetric CSM at a larger radius. The
asymmetric CSM is confirmed in a multicomponent H$\rm{\alpha}$ profile in a
nebular spectrum from day 142.

Given the peculiarity of the light curve and spectroscopic evolution of SN
2018ivc, the qualitative similarity to SN 1996al during the first 80 days of
evolution (before it disappeared behind the Sun) is striking. The light curves
decline at similar rates and show similar changes in slope, including a short
plateau. Additionally, their spectra are dominated by emission lines from the
same species. The steeply declining light curve and short plateau of SN
2018ivc could be a sign of a low mass of hydrogen in its ejecta. Benetti et
al. (2016) model the light curve of SN 1996al and find that the light from the
ejecta is significantly fainter than the light from the CSM. The fact that SN
2018ivc is less luminous than SN 1996al indicates that the majority of the
light may be coming from the SN ejecta, although the frequent changes in slope
demonstrate that CSM-ejecta interaction is significant at some phases.

Additionally, the similarity of the hydrogen and He i emission profiles
indicates that they originate in the same part of the ejecta and that perhaps,
like SN 1996al, the ejecta of SN 2018ivc are predominantly helium. While
multipeaked emission profiles are not observed in the early-time spectra of SN
2018ivc, asymmetry is seen spectroscopically in the HV feature that appears
around day 18. The broad H$\rm{\alpha}$ profile in the day 279 spectrum of SN
2018ivc does not show multiple components. It is possible the SN 2018ivc has a
multi-component H$\rm{\alpha}$ profile, like the one clearly visible in SN
1996al, but that the two components are blended to the point of being
indistinguishable from a single profile. Given the similarity between SN
2018ivc and SN 1996al, it is possible that they have similar progenitors as we
will discuss in Section 7.

###  6.2 Progenitor from HST Pre-Imaging

High-resolution images taken prior to explosion can be used to identify and
characterize the properties of a SN progenitor (e.g. Crockett et al., 2011;
Elias-Rosa et al., 2010, 2011; Fraser et al., 2010, 2011, 2014; Kochanek et
al., 2012; Li et al., 2007; Mattila et al., 2008; Maund & Smartt, 2005, 2009;
Maund et al., 2013, 2014a; Smartt et al., 2004; Smartt, 2009; Tomasella et
al., 2013b; Maund et al., 2014b; Kochanek et al., 2017; Kilpatrick & Foley,
2018; Van Dyk et al., 2019; Van Dyk, 2017). We located pre-explosion HST
observations in the Mikulski Archive for Space Telescopes (MAST) and analyzed
them for the presence of a progenitor. The SN site is located in Advanced
Camera for Surveys/Wide Field Channel (ACS/WFC) data in bands F658N and F814W
from program GO-9788 (PI: L. Ho) and in F550M from GO-9503 (PI: N. Nagar), as
well as Wide Field Planetary Camera (WFPC2) images in F606W from both GO-5479
(PI: M. Malkan) and GO-8597 (PI: M. Regan) and in F450W from GO-11128 (PI: D.
Fisher); see Table 2.

To precisely pinpoint the SN location in the archival data, we subsequently
obtained higher spatial resolution HST images of the SN with the Wide Field
Camera 3 (WFC3) on 2019 July 1. We identify 69 sources in common between
between the F814W WFC3 image of the SN to the GO-9788 F814W ACS exposure and
use these identify the SN location in the pre-explosion ACS image. To do this,
we randomly select 34 sources and use these to compute the astrometric
transformation from the WFC3 image to the ACS image the PyRAF task geomap. We
then measure the location of the SN in the pre-explosion image using the PyRAF
task geoxytran and the astrometric transformation from geomap. To understand
the error introduced by the sources used to find the astrometric
transformation, we calculate the root-mean-square (RMS) uncertainty of the 35
stars not used to determine the transformation. We repeat this process 1000
times and find the SN to be located at pixel ($\rm{3546.990\pm 0.011\pm
0.062}$, $\rm{4069.583\pm 0.008\pm 0.049}$) in the ACS F814W image, where the
first uncertainty reported is the standard deviation of the measured SN
location over the 1000 trials and the second uncertainty corresponds to the
median RMS uncertainty of the stars not used to calculate the astrometric
transformation. No source is detected within five sigma of this location, as
indicated in Figure 11.

We then processed the individual archival FLC and C0F frames through
AstroDrizzle (Hack et al., 2012) to flag cosmic-ray hits, and then extracted
photometry from these frames using Dolphot (Dolphin, 2000, 2016). Given that
the prognitor is not detected in the pre-explosion HST images, we place upper
limits on the progenitor detection which we list in Table 2.

We note in passing that, owing to the relative proximity of SN 2018ivc to the
active nucleus of NGC 1068, the various pre-explosion data of the host galaxy
obtained by the Spitzer Space Telescope, even at the shortest wavelength IRAC
band at 3.6 $\mu$m, are of little value for progenitor identification, since
the image of the nucleus was too luminous and effectively saturated the
detectors. Given the comparatively low Spitzer spatial resolution, the pixels
at the SN site were heavily affected by this saturation.

Figure 11: Left: Pre-explosion HST ACS/WFC image from 2003 Oct. 26 at F814W of NGC 1068 containing the site of SN 2018ivc. No source is identified at the SN location, indicated by the red ticks. Right: HST WFC3/UVIS image at F814W of SN 2018ivc, obtained on 2019 July 1. The SN is indicated with the red ticks. Table 2: HST Upper Limits to the SN 2018ivc Progenitor Detection Date | Instrument | Filter | Apparent Mag | Absolute Mag | Program ID | PI
---|---|---|---|---|---|---
|  |  | Limit (3$\sigma$) | Limit (3$\sigma$) |  |
2007-08-16 | WFPC2 | F450W | $>25.8$ | $>-6.2$ | GO-11128 | D. Fisher
2003-01-08 | ACS/WFC | F550M | $>25.0$ | $>-6.5$ | GO-9503 | N. Nagar
1994-12-03 | WFPC2 | F606W | $>25.0$ | $>-6.4$ | GO-5479 | M. Malkan
2001-06-30 | WFPC2 | F606W | $>25.4$ | $>-6.0$ | GO-8597 | M. Regan
2003-10-26 | ACS/WFC | F658N | $>23.4$ | $>-7.9$ | GO-9788 | L. Ho
2003-10-26 | ACS/WFC | F814W | $>24.4$ | $>-6.5$ | GO-9788 | L. Ho

##  7 Discussion

The immediate identification and high-cadence photometric and spectroscopic
follow-up observations offer us a detailed picture of the evolution of SN
2018ivc that would not otherwise be possible. Although the declining light
curve indicates that this is a Type IIL-like SN, there is evidence that the
progenitor is more complicated than that of the typical Type IIL-like SN. The
strong He i lines, not always visible in Type IIL-like SNe, could indicate
that the progenitor lost most of its hydrogen envelope. The rapidly declining
light curve corroborates this picture of mass loss. The frequent change in
slope suggests that, in addition to the linear decay of the small hydrogen
envelope, the shock is encountering shells of different densities that were
ejected from the star during its lifetime. This would imply that some
interaction between the SN and the CSM is also partially powering the light
curve of SN 2018ivc, even though narrow lines typical of some interacting SNe
(IIn) are not detected.

The narrow lines ($\sim 10^{2}$ $\rm{km~{}s^{-1}}$) of Type IIn SNe are formed
by the recombination of unshocked CSM that has been ionized by photons from
the forward-shock front. CSM can also produce intermediate-width lines ($\sim
10^{3}$ $\rm{km~{}s^{-1}}$) from the recombination of gas after the shock wave
has passed through it. In SN 2018ivc, we observe broad emission from the SN
ejecta and narrow emission from the host galaxy, but find no evidence that a
narrow line SN component exists (see Section 5.2). This lack of narrow lines
may indicate a clumpy CSM that has been enveloped by the SN ejecta (Smith et
al., 2015; Andrews & Smith, 2018). We note, however, that Type IIn SNe do not
always exhibit narrow lines at all epochs and it is possible that we do not
have a high-resolution observation during the time that narrow lines were
visible. Additionally, it is possible that the narrow lines from the CSM are
not strong enough to show above broad lines of the ejecta or are too narrow
and weak to be distinguished from the significant host contamination.

Despite the lack of narrow lines, there are several other indications that the
ejecta of SN 2018ivc are interacting with CSM. The boxy profiles of the
H$\rm{\alpha}$ and He i $\lambda 6678$ complex could also be indicative of
interaction with a shell of CSM (Andrews et al., 2010; Inserra et al., 2011).
The strong X-ray detection likely originates from the shocked CSM. Together,
the light curve and spectroscopic observations demonstrate the presence of
interaction in SN 2018ivc.

The HV features seen in SN 2018ivc are unusual. The presence of these features
indicates that hot, dense, asymmetric material is moving with the speed of the
ejecta (if the material is in the line of sight) or faster. Given that the HV
features are present in hydrogen, helium, and calcium emission, this feature
may be due to material that was ejected in the explosion. It is possible that
a bullet of $\rm{{}^{56}Ni}$ was ejected early in the explosion at high speed
and that its the radioactive decay powers these features.

Searching the literature, we find that while multicomponent hydrogen features
are seen in Type IIP/IIL SNe, they are most often during the nebular phase and
at significantly lower velocity. Two notable exceptions to this are SN 2014G
(Terreran et al., 2016) and SN 2010jp (Smith et al., 2012). SN 2014G was a
Type IIL-like SN which showed narrow flash spectroscopic features during the
first 10 days of evolution. Around day 100, SN 2014G developed a narrow
feature blueward of H$\rm{\alpha}$ that could not be associated with any other
species and thus was identified as a HV hydrogen feature, with an initial
velocity of $\sim 7600$ $\rm{km~{}s^{-1}}$. Terreran et al. (2016) explain
this feature as being caused by the spherically symmetric SN ejecta
interacting with a bipolar lobe CSM with a $40^{\circ}$ angle between the CSM
axis and the observer’s line of sight. SN 2010jp was a low luminosity,
linearly declining Type IIn SN which showed a triple-peaked H$\rm{\alpha}$
emission line. The red and blue H$\rm{\alpha}$ peaks ($-12,000$ and 15,000
$\rm{km~{}s^{-1}}$, respectively) are explained by a jet-powered explosion.
The HV features in SN 2018ivc, at comparable speeds to those of SN 2014G and
SN 2010jp, could originate from a disk or jet-like structure.

Based on the upper limits we derive from the HST pre-explosion images (see
Table 2 and Figure 12) and the fact that hydrogen is visible and strong
throughout the spectroscopic evolution, we evaluate what we can infer about
the progenitor. Referring to the single-star evolutionary tracks at solar
metallicity from the MESA Isochrones & Stellar Tracks (MIST v1.2; Choi et al.
2016; Dotter 2016; Paxton et al. 2011, 2013, 2015), we exclude stars whose
photometry would exceed our upper limits or with less than 0.1
$\rm{M_{\odot}}$ of hydrogen in the envelope. With these criteria we find we
can eliminate all stars with initial masses between 9 $\rm{M_{\odot}}$ and 48
$\rm{M_{\odot}}$ and above 52 $\rm{M_{\odot}}$ as the progenitor (see left
panel of Figure 12). We conservatively include stars between 49
$\rm{M_{\odot}}$ and 52 $\rm{M_{\odot}}$ although we note that the largest
hydrogen envelope mass in this range is 0.5 $\rm{M_{\odot}}$, which could
possibly be excluded with detailed spectroscopic modeling that is beyond the
scope of this paper. Additionally, we caution that this limit is highly
dependent on the evolutionary models used. For instance, if we adopt the STARS
models (Eggleton 1971; Pols et al. 1995; Eldridge & Tout 2004; on which the
BPASS binary evolution models are based, e.g., Eldridge et al. 2017), we find
that the hydrogen mass in the envelope drops below 0.1$\rm{M_{\odot}}$ at a
lower mass and redder SED, resulting in the exclusion of all stars with masses
greater than 8 $\rm{M_{\odot}}$ (see right panel of Figure 12).

Figure 12: Upper limits on the detection, as given in Table 2, of the SN
2018ivc progenitor in archival HST images (black squares and arrows). In color
are model SEDs derived from the endpoint of the solar-metallicity MIST (Choi
et al., 2016; Dotter, 2016; Paxton et al., 2011, 2013, 2015) single-star
evolutionary tracks (left) and the STARS models (Eggleton 1971; Pols et al.
1995; Eldridge & Tout 2004 (right) for masses spanning the range of allowed
masses for each evolutionary model code. Allowed models are denoted with
filled color circles. Models that are ruled out by the upper limits are marked
with open circles and connected by a dashed line. Models that are ruled out by
a lack of hydrogen in their envelope ($M\rm{{}_{H,env}}<0.1$ $\rm{M_{\odot}}$)
are denoted with open circles and connected by a dotted line. The HST upper
limits and strong hydrogen features in the spectra imply that that progenitor
of SN 2018ivc, if single, was likely 8 $\rm{M_{\odot}}$. The MIST models allow
additional progenitors in the range 49-52 $\rm{M_{\odot}}$, although these
have small hydrogen envelopes which may not be able to produce the features
seen in the spectra of SN 2018ivc.

The majority of massive stars form in binary systems (Sana et al., 2012). For
this reason, we also consider possible binary progenitor systems by examining
the endpoints of the BPASS v2.2 models and the light from each combined
system. Again, considering the HST upper limits and the mass of hydrogen in
the envelope, we find a maximum progenitor mass of 11 $\rm{M_{\odot}}$, with
the majority of progenitors between 8 $\rm{M_{\odot}}$ and 9 $\rm{M_{\odot}}$
(see Figure 13). Most progenitor systems are in a wide binary with
$\rm{log(Period[day])=2.75}$ for masses between 8 $\rm{M_{\odot}}$ and 11
$\rm{M_{\odot}}$ and a broader range of $\rm{log(Period)}$ (2.5-4) for an 8
$\rm{M_{\odot}}$ progenitor. The few short-period progenitors occur when the
secondary mass is much smaller than the primary mass ($\rm{M_{2}/M_{1}\leq
0.2}$; see Figure 14).

Figure 13: The range of progenitor masses and periods of the solar metallicity
model binary systems from BPASS (Eldridge et al., 2017) that are allowed by
the pre-explosion upper limits. Models that also contain more than 0.1
$\rm{M_{\odot}}$ of hydrogen in their envelopes are marked as black circles,
while models that contain less hydrogen are marked as pink points. The model
parameters are marginalized over the ratio of the secondary to primary star
mass (see Figure 14 for an animation of the three-dimensional distribution).
The opacity of the markers indicates the density of models at a given
location. The distribution of progenitor masses, further marginalized over the
period, is shown in the top panel, with allowed progenitors in black and
eliminated progenitors in pink. Similarly, the distribution of periods,
further marginalized over progenitor mass, is shown in the right panel with
the same color scheme. Consistent with the single stars, we find that the
progenitor was less than 12 $\rm{M_{\odot}}$ and most probably 8
$\rm{M_{\odot}}$. Figure 14: An animated version of this figure is available
in the online version of this paper. The range of progenitor mass, period, and
mass ratio of the secondary to the primary star of the solar metallicity model
binary systems from BPASS (Eldridge et al., 2017) that are allowed by the pre-
explosion upper limits. Models that are allowed by both the HST upper limits
and the hydrogen in their envelopes are shown in black, models that are ruled
out are shown in pink. Most mass ratios are possible for allowed
$\rm{log(Period[days])}$ and progenitor masses, although there are 6 models
that only occur at specific combinations of all three parameters.

To better understand the nature of SN 2018ivc and its progenitor, we compare
the host properties derived in Section 4.2 with those of all Type IIP/IIL SNe
(85) and Type IIn SNe (16) from the PMAS/PPak Integral-field Supernova hosts
COmpilation (PISCO) sample (Galbany et al., 2018b)777Updated with all new
observations obtained through May 2019.. The metallicity of the parent H ii
region of SN 2018ivc is near the median of the Type IIP/IIL SN distribution
and slightly below the median of the Type IIn SN distribution (see the upper-
right panel of Figure 6). Similarly, the EWH$\rm{\alpha}$ falls near the
median of both Type IIP/IIL SN and Type IIn SN distributions.

EWH$\rm{\alpha}$ is an indicator of the age of the cluster. Kuncarayakti et
al. (2013) show a theoretical relationship between EWH$\rm{\alpha}$ and the
age of the stellar population assuming a Salpeter initial mass function (IMF)
and an instantaneous burst of star formation. The age of the stellar
population can then be used to estimate the age of the SN progenitor. Using
the EWH$\rm{\alpha}$ derived from the parent H ii region and Figure 1 from
Kuncarayakti et al. (2013), we find the age of the progenitor of SN 2018ivc to
be 6.75–7.75 Myr (depending on the slope and upper value of the IMF),
corresponding to a progenitor mass of 25–28 M⊙. This value is similar to the
25 M⊙ of SN 1996al found by Benetti et al. (2016). We note, however, that this
agreement may be coincidental owing to limitations in our ability to isolate a
single stellar population at this distance as well as simplifying assumptions
made in mapping of the EWH$\rm{\alpha}$ modeling to progenitor age (see Schady
et al. 2019 for a detailed explanation). Specifically, the
EWH$\rm{\alpha}$-age relation used assumes that massive stars are the ionizing
source producing the H$\rm{\alpha}$ emission, that we recover all of the
photons ionizing the surrounding gas (i.e., there is no leakage), and that
there are no binaries in the stellar population. The inclusion of these
effects could yield an older age and therefore a lower mass progenitor, which
is more consistent with the constraints from pre-explosion imaging.

##  8 Summary

In this paper we have described the early discovery and prompt follow-up
observations of SN 2018ivc by the DLT40 team and high-cadence monitoring by
the GSP over the first 80 days of evolution. The DLT40 survey observed SN
2018ivc, identified it as a SN candidate using automated software, and
confirmed the SN with follow-up observations within 15 min of the discovery
observation. This discovery triggered a comprehensive set of photometric and
spectroscopic observations by the DLT40 team and the GSP for the duration of
visibility, $\sim 80$ days.

The light curve of this SN changed slope every $\approx 10$ days for the first
30 days before settling onto a linear decline. The spectroscopic evolution is
rapid and the spectrum is dominated by H and He i emission lines with shallow
or no P Cygni absorption. These characteristics combined with the X-ray
detection suggest that the SN ejecta are interacting with CSM. We find HV
emission in hydrogen, helium, and calcium, indicating an asymmetric progenitor
or explosion. We analyze pre-explosion IFU observations of NGC 1068 and find
that SN 2018ivc exploded in a region of typical metallicity and star formation
with respect to other Type II SNe. By analyzing the EWH$\rm{\alpha}$ of the
SN’s parent H ii region, we find evidence that the progenitor of SN 2018ivc
had an initial mass of $\rm{\leq 25\ M_{\odot}}$.

We show that SN 2018ivc resembles SN 1996al both photometrically and
spectroscopically, indicating that the progenitor of SN 2018ivc may have been
similar to the 25 $\rm{M_{\odot}}$ progenitor inferred for SN 1996al, although
this is in tension with the progenitors derived from pre-explosion
observations. Finally, we use HST archival observations of NGC 1068, taken
prior to explosion and ToO HST observations of the SN to derive photometric
upper limits on the progenitor. From these limits and the strong presence of
hydrogen in the spectra, we infer a probable low-mass progenitor ($M<12$
$\rm{M_{\odot}}$ for binary models and $M<8$ $\rm{M_{\odot}}$ for single-star
models) although the MIST models do allow for a massive progenitor ($49\leq
M\leq 52$ $\rm{M_{\odot}}$).

It is only with the early discovery, immediate response, and high-cadence
monitoring of SNe, like those provided by the DLT40 and GSP teams, that the
complex and rapid evolution of interacting SNe like SN 2018ivc can be observed
and characterized. It is possible that the interaction in other Type IIL-like
SNe, lacking narrow lines and without early identification and sufficient
follow-up observations, has not been identified. For a clear and detailed
understanding of the progenitor systems of Type II SNe, specifically the role
of mass loss in stellar evolution and its observational signatures, we must
continue to identify SNe early, announce the discovery immediately, and
combine our resources for high-cadence observations of the full SN evolution.

We are grateful to B. J. Fulton for providing us the image shown in Figure 1,
to P. Chandra for analysis of the Swift XRT and GMRT data, and to Jared
Goldberg, Charlie Kilpatrick, Morgan Fraser, and Stephen Smartt for insightful
conversations. Research by K.A.B, S.V., and Y.D. is supported by NSF grant
AST–1813176. Research by D.J.S. is supported by NSF grants AST–1821967,
1821987, 1813708, 1813466, and 1908972. Research by J.E.A. and N.S. was
supported by NSF grant AST–151559. L.G. was funded by the European Union’s
Horizon 2020 research and innovation programme under the Marie Skłodowska-
Curie grant agreement No. 839090. J.B., D.H., C.M., and D.A.H. are supported
by NSF grant AST–1313484. I.A. is a CIFAR Azrieli Global Scholar in the
Gravity and the Extreme Universe Program and acknowledges support from that
program, from the Israel Science Foundation (grant numbers 2108/18 and
2752/19), from the United States - Israel Binational Science Foundation (BSF),
and from the Israeli Council for Higher Education Alon Fellowship. A.V.F. is
grateful for financial assistance from the Christopher R. Redlich Fund, the
TABASGO Foundation, and the Miller Institute for Basic Research in Science
(U.C. Berkeley). The research of J.C.W. is supported by NSF grant AST–1813825.
K.M. acknowledges the support from Department of Science and Technology (DST),
Government of India and Indo-US Science and Technology Forum (IUSSTF) for the
WISTEMM fellowship and the Department of Physics, UC Davis, where part of this
work was carried out. P.J.B.’s work on SOUSA and core-collapse supernovae is
supported by NASA ADAP grants NNX13AF35G and NNX17AF43G. D.P. gratefully
acknowledges support provided by the National Aeronautics and Space
Administration (NASA) through Chandra Award Number GO8-19051X issued by the
Chandra X-ray Center, which is operated by the Smithsonian Astrophysical
Observatory for and on behalf of NASA under contract NAS8-03060. J.Z. is
supported by the National Natural Science Foundation of China (NSFC, grants
11773067 and 11403096), the Youth Innovation Promotion Association of the CAS
(grant 2018081), and the Western Light Youth Project. This work of X.W. is
supported by the National Natural Science Foundation of China (NSFC grants
11325313 and 11633002), and the National Program on Key Research and
Development Project (grant 2016YFA0400803). E.B. and J.D. are supported in
part by NASA grant NNX16AB25G. This work was partially supported by the Open
Project Program of the Key Laboratory of Optical Astronomy, National
Astronomical Observatories, Chinese Academy of Sciences. We thank the support
staffs at the many observatories where data were obtained. Funding for the Li-
Jiang 2.4 m telescope (LJT) has been provided by Chinese Academy of Sciences
and the People’s Government of Yunnan Province. The LJT is jointly operated
and administrated by Yunnan Observatories and Center for Astronomical Mega-
Science, CAS. This work makes use of observations from the Las Cumbres
Observatory network. Based in part on observations obtained at the Gemini
Observatory under programs GS-2018B-Q-130 (PI: D. Sand). Gemini is operated by
the Association of Universities for Research in Astronomy, Inc., under a
cooperative agreement with the NSF on behalf of the Gemini partnership: the
NSF (United States), the National Research Council (Canada), CONICYT (Chile),
Ministerio de Ciencia, Tecnología e Innovación Productiva (Argentina), and
Ministério da Ciência, Tecnologia e Inovação (Brazil). The data were processed
using the Gemini IRAF package. We thank the queue service observers and
technical support staff at Gemini Observatory for their assistance. This work
is based in part on observations made with the NASA/ESA Hubble Space
Telescope, obtained from the Data Archive at the Space Telescope Science
Institute (STScI), which is operated by the Association of Universities for
Research in Astronomy (AURA), Inc., under NASA contract NAS 5-26555. Support
for program GO-15151 was provided by NASA through a grant from STScI. 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 research has made use of the NASA/IPAC Extragalactic Database (NED) which
is operated by the Jet Propulsion Laboratory, California Institute of
Technology, under contract with NASA. Observations using Steward Observatory
facilities were obtained as part of the large observing programme AZTEC:
Arizona Transient Exploration and Characterization. Some observations reported
here were obtained at the MMT Observatory, a joint facility of the University
of Arizona and the Smithsonian Institution. The LBT/MODS observations were
obtained as part of a pilot effort at Steward Observatory to move to queue
observing and include back-up programs for poor weather conditions or
technical failures. We thank J. Power for coordinating the queue efforts for
our LBT/MODS observations. Based in part on observations obtained at the
Southern Astrophysical Research (SOAR) telescope, which is a joint project of
the Ministério da Ciência, Tecnologia, Inovações e Comunicações (MCTIC) do
Brasil, the U.S. National Optical Astronomy Observatory (NOAO), the University
of North Carolina at Chapel Hill (UNC), and Michigan State University (MSU).
The facilities at IAO and CREST are operated by the Indian Institute of
Astrophysics, Bangalore. This research has made use of data obtained from the
Chandra Data Archive and software provided by the Chandra X-ray Center (CXC)
in the application packages CIAO and Sherpa. Some of the data presented herein
were obtained at the W. M. Keck Observatory, which is operated as a scientific
partnership among the California Institute of Technology, the University of
California, and NASA; the observatory was made possible by the generous
financial support of the W. M. Keck Foundation. , Spitzer (IRAC, MIPS), Swift
(UVOT), YAO:2.4m (YFOSC)

##  Appendix A Spectroscopic Observations

Table 3: Log of Spectroscopic Observations Observation | MJD | Phase | Telescope | Instrument | Wavelength | Exposure | Resolution
---|---|---|---|---|---|---|---
Date |  | (day) |  |  | Range (Å) | Time (s) | $\rm{\lambda/\Delta\lambda}$
2018-11-24 | 58446.29 | 2.04 | FTN | FLOYDS | 3499-10000 | 2700 | 400-700
2018-11-24 | 58446.50 | 2.25 | 2.16m XLO | BFOSC | 3852-8696 | 3000 | 250-800
2018-11-24 | 58446.70 | 2.45 | 2.4m LJT | YFOSC | 3400-9100 | 1200 | 240
2018-11-24 | 58446.81 | 2.56 | SALT | RSS | 3494-9392 | 1793 | 600-2000
2018-11-25 | 58447.20 | 2.95 | Gemini-S | FLAMINGOS-2 | 9852-18075 | 720 | 900
2018-11-25 | 58447.38 | 3.13 | FTN | FLOYDS | 3499-10000 | 1200 | 400-700
2018-11-25 | 58447.52 | 3.27 | 2.16m XLO | BFOSC | 3852-8696 | 3300 | 250-800
2018-11-25 | 58447.60 | 3.35 | HCT-IIA | HFOSC | 3500-8998 | 1200 | 1200
2018-11-26 | 58448.57 | 4.32 | 2.16m XLO | BFOSC | 3852-8697 | 3300 | 250-800
2018-11-27 | 58449.16 | 4.91 | HET | lrs2 | 3640-10298 | 606 | 1140-1920
2018-11-27 | 58449.32 | 5.07 | LBT | MODS | 3290-5549 | 2400 | 1850-2300
2018-11-27 | 58449.32 | 5.07 | LBT | MODS | 5800-9572 | 2400 | 1850-2300
2018-11-27 | 58449.61 | 5.36 | 2.16m XLO | BFOSC | 3853-8701 | 3300 | 250-800
2018-11-28 | 58450.23 | 5.98 | FTN | FLOYDS | 3499-10000 | 1200 | 400-700
2018-11-29 | 58451.25 | 7.00 | Bok | SPOL | 4001-7549 | 2400 | 430
2018-12-01 | 58453.40 | 9.15 | FTN | FLOYDS | 3499-10000 | 1200 | 400-700
2018-12-02 | 58454.00 | 9.75 | Gemini-S | FLAMINGOS-2 | 9851-18077 | 1080 | 900
2018-12-03 | 58455.69 | 11.44 | HCT-IIA | HFOSC | 3500-8998 | 1200 | 1200
2018-12-04 | 58456.59 | 12.34 | HCT-IIA | HFOSC | 3500-8998 | 1200 | 1200
2018-12-05 | 58457.44 | 13.19 | Keck I | LRIS | 3136-10220 | 119 | 750-1475
2018-12-07 | 58459.20 | 14.95 | Gemini-S | FLAMINGOS-2 | 9849-18077 | 1080 | 900
2018-12-07 | 58459.55 | 15.30 | HCT-IIA | HFOSC | 3500-8998 | 1800 | 1200
2018-12-08 | 58460.36 | 16.11 | FTN | FLOYDS | 3500-9999 | 1200 | 400-700
2018-12-10 | 58462.32 | 18.07 | 3.5m APO | DIS | 3374-5607 | 2400 | 2450
2018-12-10 | 58462.32 | 18.07 | 3.5m APO | DIS | 5263-9404 | 2400 | 3150
2018-12-11 | 58463.81 | 19.56 | HCT-IIA | HFOSC | 3500-8998 | 2400 | 1200
2018-12-13 | 58465.00 | 20.75 | Gemini-S | FLAMINGOS-2 | 9850-18077 | 1440 | 900
2018-12-14 | 58466.30 | 22.05 | MMT | Binospec | 5072-7541 | 720 | 3590
2018-12-15 | 58467.76 | 23.51 | HCT-IIA | HFOSC | 3500-8998 | 2400 | 1200
2018-12-16 | 58468.65 | 24.40 | HCT-IIA | HFOSC | 3500-8998 | 2400 | 1200
2018-12-17 | 58469.23 | 24.98 | FTN | FLOYDS | 3500-10000 | 1800 | 400-700
2018-12-18 | 58470.51 | 26.26 | 2.16m XLO | BFOSC | 4358-8690 | 2400 | 250-800
2018-12-19 | 58471.10 | 26.85 | Gemini-S | FLAMINGOS-2 | 9852-18076 | 1800 | 900
2018-12-21 | 58473.15 | 28.90 | SOAR | GHTS | 3500-7000 | 780 | 1050
2018-12-21 | 58473.15 | 28.90 | SOAR | GHTS | 5000-9000 | 780 | 1050
2018-12-22 | 58474.35 | 30.10 | FTN | FLOYDS | 3499-9999 | 3600 | 400-700
2018-12-23 | 58475.21 | 30.96 | FTN | FLOYDS | 3499-9999 | 3600 | 400-700
2018-12-24 | 58476.63 | 32.38 | HCT-IIA | HFOSC | 3500-8998 | 2400 | 1200
2018-12-24 | 58476.78 | 32.53 | HCT-IIA | HFOSC | 3500-8998 | 1200 | 1200
2018-12-24 | 58476.85 | 32.60 | SALT | RSS | 6057-7010 | 2464 | 2200-5500
2018-12-26 | 58478.10 | 33.85 | Gemini-S | FLAMINGOS-2 | 9849-18077 | 1800 | 900
2018-12-26 | 58478.48 | 34.23 | FTS | FLOYDS | 4800-10000 | 3600 | 400-700
2018-12-26 | 58478.83 | 34.58 | SALT | RSS | 3497-9393 | 2093 | 600-2000
2018-12-30 | 58482.49 | 38.24 | FTS | FLOYDS | 4800-9999 | 3600 | 400-700
2019-01-03 | 58486.30 | 42.05 | FTN | FLOYDS | 3500-10000 | 3600 | 400-700
2019-01-05 | 58488.00 | 43.75 | Gemini-S | FLAMINGOS-2 | 9851-18075 | 1800 | 900
2019-01-06 | 58489.54 | 45.29 | 2.16m XLO | BFOSC | 3851-8692 | 3600 | 250-800
2019-01-10 | 58493.33 | 49.08 | FTN | FLOYDS | 3499-10000 | 3600 | 400-700
2019-01-11 | 58494.30 | 50.05 | Keck I | LRIS | 3137-10197 | 420 | 750-1475
2019-01-13 | 58496.49 | 52.24 | 2.16m XLO | BFOSC | 4095-8810 | 3000 | 250-800
2019-01-14 | 58497.10 | 52.85 | Gemini-S | FLAMINGOS-2 | 9850-18079 | 2160 | 900
2019-01-22 | 58505.26 | 61.01 | FTN | FLOYDS | 3499-9999 | 3600 | 400-700
2019-01-28 | 58511.44 | 67.19 | 2.16m XLO | BFOSC | 3869-8822 | 3600 | 250-800
2019-02-05 | 58519.25 | 75.00 | Keck I | LRIS | 3138-10244 | 1200 | 750-1475
2019-02-08 | 58522.13 | 77.88 | MMT | Binospec | 5062-7522 | 960 | 3590
2019-08-28 | 58723.59 | 279.34 | Keck II | DEIMOS | 4480-9510 | 1200 | 1875

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