# ASASSN-15hy: An Underluminous, red 03fg-like Type Ia Supernova

J. Lu (陆晶) Department of Physics, Florida State University, 77 Chieftan Way,
Tallahassee, FL 32306, USA  C. Ashall Institute for Astronomy, University of
Hawaii, 2680 Woodlawn Drive, Honolulu, HI 96822,USA  E. Y. Hsiao Department of
Physics, Florida State University, 77 Chieftan Way, Tallahassee, FL 32306, USA
P. Hoeflich Department of Physics, Florida State University, 77 Chieftan Way,
Tallahassee, FL 32306, USA  L. Galbany Institute of Space Sciences (ICE,
CSIC), Campus UAB, Carrer de Can Magrans, s/n, E-08193 Barcelona, Spain  E.
Baron Homer L. Dodge Department of Physics and Astronomy, University of
Oklahoma, Rm 100 440 W. Brooks, Norman, OK 73019-2061, USA  Hamburger
Sternwarte, Gojenbergsweg 112, D-21029 Hamburg, Germany  M. M. Phillips Las
Campanas Observatory, Carnegie Observatories, Casilla 601, La Serena, Chile
C. Contreras Las Campanas Observatory, Carnegie Observatories, Casilla 601, La
Serena, Chile  C. R. Burns Observatories of the Carnegie Institution for
Science, 813 Santa Barbara Street, Pasadena, CA 91101, USA  N. B. Suntzeff
George P. and Cynthia Woods Mitchell Institute for Fundamental Physics and
Astronomy, Department of Physics and Astronomy, Texas A&M University, College
Station, TX 77843, USA  M. D. Stritzinger Department of Physics and Astronomy,
Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark.  J. Anais Las
Campanas Observatory, Carnegie Observatories, Casilla 601, La Serena, Chile
J. P. Anderson European Southern Observatory, Alonso de Córdova 3107, Casilla
19, Santiago, Chile  P. J. Brown George P. and Cynthia Woods Mitchell
Institute for Fundamental Physics and Astronomy, Department of Physics and
Astronomy, Texas A&M University, College Station, TX 77843, USA  L. Busta Las
Campanas Observatory, Carnegie Observatories, Casilla 601, La Serena, Chile
S. Castellón  Las Campanas Observatory, Carnegie Observatories, Casilla 601,
La Serena, Chile  S. Davis Department of Physics and astronomy, University of
California, 1 Shields Avenue, Davis, CA 95616-5270, USA  T. Diamond Department
of Physics, Florida State University, 77 Chieftan Way, Tallahassee, FL 32306,
USA  E. Falco Harvard-Smithsonian Center for Astrophysics, 60 Garden Street,
Cambridge, MA 02138, USA  C. Gonzalez  Las Campanas Observatory, Carnegie
Observatories, Casilla 601, La Serena, Chile  M. Hamuy Vice President and Head
of Mission of AURA-O in Chile, Avda Presidente Riesco 5335 Suite 507,
Santiago, Chile  Hagler Institute for Advanced Studies, Texas A&M University,
College Station, TX 77843, USA  S. Holmbo Department of Physics and Astronomy,
Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark.  T. W.-S.
Holoien Observatories of the Carnegie Institution for Science, 813 Santa
Barbara Street, Pasadena, CA 91101, USA  K. Krisciunas George P. and Cynthia
Woods Mitchell Institute for Fundamental Physics and Astronomy, Department of
Physics and Astronomy, Texas A&M University, College Station, TX 77843, USA
R. P. Kirshner Harvard-Smithsonian Center for Astrophysics, 60 Garden Street,
Cambridge, MA 02138, USA  Gordon and Betty Moore Foundation, 1661 Page Mill
Road, Palo Alto, CA 94304, USA  S. Kumar Department of Physics, Florida State
University, 77 Chieftan Way, Tallahassee, FL 32306, USA  H. Kuncarayakti
Tuorla Observatory, Department of Physics and Astronomy, FI-20014 University
of Turku, Finland  Finnish Centre for Astronomy with ESO (FINCA), FI-20014
University of Turku, Finland  G. H. Marion  University of Texas at Austin, 1
University Station C1400, Austin, TX 78712-0259, USA  N. Morrell Las Campanas
Observatory, Carnegie Observatories, Casilla 601, La Serena, Chile  S. E.
Persson Observatories of the Carnegie Institution for Science, 813 Santa
Barbara St, Pasadena, CA 91101, USA  A. L. Piro Observatories of the Carnegie
Institution for Science, 813 Santa Barbara St, Pasadena, CA 91101, USA  J. L.
Prieto  Nucleo de Astronomía de la Facultad de Ingeniería y Ciencias,
Universidad Diego Portales, Av. Ejército 441, Santiago, Chile  Millennium
Institute of Astrophysics, Santiago, Chile  D. J. Sand Steward Observatory,
University of Arizona, 933 North Cherry Avenue, Rm. N204, Tucson, AZ
85721-0065, USA  M. Shahbandeh Department of Physics, Florida State
University, 77 Chieftan Way, Tallahassee, FL 32306, USA  B. J. Shappee
Institute for Astronomy, University of Hawaii, 2680 Woodlawn Drive, Honolulu,
HI 96822,USA  F. Taddia Department of Physics and Astronomy, Aarhus
University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark.

(Received March 29, 2021; Accepted July 16,2021)

###### Abstract

We present photometric and spectroscopic observations of the 03fg-like type Ia
supernova (SN Ia) ASASSN-15hy from the ultraviolet (UV) to the near-infrared
(NIR). ASASSN-15hy shares many of the hallmark characteristics of 03fg-like
SNe Ia, previously referred to as “super-Chandrasekhar” SNe Ia. It is bright
in the UV and NIR, lacks a clear $i$-band secondary maximum, shows a strong
and persistent C ii feature, and has a low Si ii $\lambda$6355 velocity.
However, some of its properties are also extreme among the subgroup.
ASASSN-15hy is underluminous ($M_{B,\rm{peak}}=$$-19.14^{+0.11}_{-0.16}$ mag),
red ($(B-V)_{Bmax}=0.18^{+0.01}_{-0.03}$ mag), yet slowly declining
($\Delta{\rm{m}_{15}}(B)$=0.72 $\pm$ 0.04 mag). It has the most delayed onset
of the $i$-band maximum of any 03fg-like SN. ASASSN-15hy lacks the prominent
$H$-band break emission feature that is typically present during the first
month past maximum in normal SNe Ia. Such events may be a potential problem
for high-redshift SN Ia cosmology. ASASSN-15hy may be explained in the context
of an explosion of a degenerate core inside a nondegenerate envelope. The
explosion impacting the nondegenerate envelope with a large mass provides
additional luminosity and low ejecta velocities. An initial deflagration
burning phase is critical in reproducing the low 56Ni mass and luminosity,
while the large core mass is essential in providing the large diffusion time
scales required to produce the broad light curves. The model consists of a
rapidly rotating 1.47 $M_{\odot}$ degenerate core and a 0.8 $M_{\odot}$
nondegenerate envelope. This “deflagration core-degenerate” scenario may
result from the merger between a white dwarf and the degenerate core of an
asymptotic giant branch star.

supernovae: general – supernovae: Type Ia supernova – supernovae: individual
(ASASSN-15hy)

††journal: ApJ

##  1 Introduction

Type Ia supernovae (SNe Ia) are one of the most precise extragalactic distance
indicators in the cosmos. Their high intrinsic luminosity and apparent
uniformity enable them to be used in cosmological studies, such as mapping the
expansion history of the universe (e.g., Riess et al., 1998; Perlmutter et
al., 1999), measuring the Hubble constant, and determining the equation of
state of dark energy (e.g., Riess et al., 2007; Sullivan et al., 2011; Suzuki
et al., 2012).

SNe Ia are thought to come from the thermonuclear explosions of at least one
carbon-oxygen white dwarf (C/O WD) in a binary system (Hoyle & Fowler, 1960).
The explosive products are mainly dictated by nuclear physics, and, to first
order, it is the effectiveness of this burning that determines the observed
photometric and spectroscopic properties (e.g., Hoeflich & Khokhlov, 1996;
Mazzali et al., 2007). These luminous events are powered by the decay of
radioactive 56Ni. The amount of 56Ni produced determines the peak luminosity
(Arnett, 1982), while the interplay between 56Ni, luminosity, and opacity
determines the light-curve shape (Mazzali et al., 2001). For normal SNe Ia,
brighter objects produce more 56Ni and also have broader light curves
(Phillips, 1993; Nugent et al., 1995; Phillips et al., 1999).

However, the ejecta mass and the exact details of the explosions are still
debated, and the progenitor systems and explosion mechanisms of SNe Ia are
still not fully understood (e.g., Maoz et al. 2014). Some of the possible
progenitor scenarios are single degenerate, double degenerate,
triple/quadruple systems, and core degenerate. The single-degenerate scenario
consists of a C/O WD and a nondegenerate companion star which is either a red-
giant, main sequence or He star (e.g., Whelan & Iben 1973; Nomoto et al.
1984), whereas the double degenerate scenario consists of two WDs that merge
on dynamical or secular time scales (Iben & Tutukov, 1984; Webbink, 1984).
Another scenario consists of collisions of two WDs (e.g., Rosswog et al.,
2009). Based on secular Lidov-Kozai mechanism, Thompson (2011) suggested that
“something akin to a collision” may occur in a triple system. Going beyond the
secular approximations, Katz & Dong (2012),Kushnir et al. (2013),Pejcha et al.
(2013) found that direct collision between two WDs may result in a
thermonuclear explosion in a triple/quadruple system. Finally, the core-
degenerate (CD) scenario describes the merger of a WD and the degenerate core
of an asymptotic giant branch (AGB) star within its nondegenerate envelope
(Livio & Riess, 2003; Kashi & Soker, 2011; Ilkov & Soker, 2012). In each
scenario, the mass of the exploding WD may be near, below, or above the
Chandrasekhar mass limit ($M_{\text{Ch}}$) depending on the exact
configuration of the system or the type of the companion star. Furthermore,
the flame front may propagate as a deflagration, detonation or may transition
between these two propagation speeds.

As observations of SNe have accumulated, many peculiar subtypes of SNe Ia with
unique and extreme observational properties have emerged. It is currently
unclear which combinations of progenitor systems and explosion mechanisms lead
to the observed SN Ia diversity. One of the most rare peculiar subtypes is the
“super-Chandrasekhar” or 03fg-like SNe.111In this work we follow the
nomenclature of naming a subtype after the first SN of its kind discovered
(i.e., 03fg-like), as it is not clear if the ejecta mass of these objects in
fact exceeds $M_{\text{Ch}}$. To date there have only been a handful of well-
studied SNe Ia identified as members of this subtype, including: SN 2003fg
(Howell et al., 2006), SN 2006gz (Hicken et al., 2007; Maeda et al., 2009), SN
2007if (Scalzo et al., 2010; Yuan et al., 2010; Childress et al., 2011), SN
2009dc (Yamanaka et al., 2009; Tanaka et al., 2010; Silverman et al., 2011;
Taubenberger et al., 2011; Hachinger et al., 2012), SN 2012dn (Chakradhari et
al., 2014; Parrent et al., 2016; Nagao et al., 2017; Taubenberger et al.,
2019), ASASSN-15pz (Chen et al., 2019), and LSQ14fmg (Hsiao et al., 2020).
Furthermore, Ashall et al. (2021) presents data of 5 more 03fg-like SNe.

Generally, 03fg-like SNe have high peak luminosities (M${}_{B}\leq-19.3$ mag),
slow decline rates ($\Delta{\rm{m}_{15}}(B)$ $\leq 1.1$ mag), and long rise
times (typically $>$ 21 d, e.g., Silverman et al. 2011; Taubenberger et al.
2011). They also have a primary $i$-band maximum that peaks after that of the
$B$ band (e.g., Ashall et al. 2020). These objects have also been shown to be
bright in the ultraviolet (UV) in comparison to normal SNe Ia (Silverman et
al., 2011; Taubenberger et al., 2011; Brown et al., 2014b; Chen et al., 2019).
However, not all members of this subtype are overluminous. For example, SN
2012dn has a relatively fainter peak luminosity (M${}_{B}=-19.3$ mag,
Taubenberger et al. 2019) comparing to those SNe Ia that have similar decline
rate, yet it still shares the other similarities of the 03fg-like group
mentioned above.

03fg-like SNe exhibit low expansion velocity gradients and have weaker Ca ii
features than normal SNe Ia at early times (e.g., Scalzo et al. 2010;
Taubenberger et al. 2011; Chakradhari et al. 2014). Furthermore, the vast
majority of these objects show very strong C ii features that persist past
maximum light. Carbon lines, if detected in a normal SN Ia, are usually only
observed in spectra within a few days of the explosion. Another peculiar
spectroscopic feature of 03fg-like SNe is the absence of the $H$-band break
near 1.5 $\mu$m within the first month past maximum light (Taubenberger et
al., 2011; Hsiao et al., 2019). In normal SN Ia, the strength and velocity of
the $H$-band break is directly linked to the amount and distribution of 56Ni
in the ejecta (Wheeler et al., 1998; Hsiao et al., 2013; Ashall et al., 2019a,
b).

The origins of these 03fg-like events are still under debate. Most recently, a
class of interaction scenarios has emerged among the possible theoretical
interpretations. Early suggestions of single-degenerate, rapid rotating
super-$M_{\text{Ch}}$ WD (Howell et al., 2006; Hachisu et al., 2012) have been
challenged since the model simulations could not reproduce the slow expansion
velocities on top of the high luminosity regardless of the explosion
mechanism, such as pure detonation, pure deflagration and delayed detonation
(e.g., Pfannes et al., 2010; Hachinger et al., 2012; Fink et al., 2018). The
“interaction” (e.g., Yuan et al., 2010; Hachinger et al., 2012; Taubenberger
et al., 2013; Noebauer et al., 2016; Taubenberger et al., 2019) or “envelope”
models (e.g., Khokhlov et al., 1993; Hoeflich & Khokhlov, 1996; Scalzo et al.,
2010, 2014; Hsiao et al., 2020), in which the interaction between the
exploding WD and surrounding circumstellar medium (CSM) provides an additional
energy source and slows down the expansion, present a promising route to
explain these peculiar events. Recent works have shown photometric evidence of
CSM such as the NIR excess in SN 2012dn (Yamanaka et al., 2016; Nagao et al.,
2017) and possible wind interaction in LSQ14fmg (Hsiao et al., 2020). However,
the nature of the “envelope” is still under debate (e.g., pre-explosion wind
material, the product of double degenerate merger, or the envelope of an AGB
star in the core degenerate scenario; Kashi & Soker, 2011; Ilkov & Soker,
2012; Aznar-Siguán et al., 2015) and there is currently no spectral evidence
of interaction observed so far.

Extensive follow-up observations of 03fg-like SNe are rare. The advent of
wide-field and untargeted searches created an opportunity for advancement in
the studies of 03fg-like SNe, as they are generally located in low-mass, low-
luminosity environments (Howell et al., 2006; Childress et al., 2011;
Taubenberger et al., 2011; Chakradhari et al., 2014; Hsiao et al., 2020), and
therefore may have been missed by previous surveys. The Carnegie Supernova
Project-II (CSP-II; 2011–2015) obtained high-precision, rapid-cadence
observations of over 100 SNe Ia (Phillips et al., 2019; Hsiao et al., 2019),
mainly from untargeted searches. One of the objects observed by the CSP-II was
the 03fg-like SN Ia ASASSN-15hy located in a nearby low-luminosity host
galaxy.

We present photometric and spectroscopic follow-up observations of
ASASSN-15hy, a new 03fg-like object sharing the vast majority of the general
properties as previous members but also exhibiting extremeness, such as being
the faintest, reddest and having the most delayed $i$-band peak time within
the subgroup. The observations and data reductions are summarized in Section
2. The host properties and extinction are presented in Section 3. In Sections
4 and 5, the photometric and spectroscopic properties of ASASSN-15hy are
analyzed, respectively. Model simulation and interpretation are discussed in
Sections 6 and 6.3. Finally, we conclude in Section 7.

##  2 Observations

Figure 1: The observed UV, optical, and NIR light curves of ASASSN-15hy. The solid markers represent CSP-II photometry. $Swift$ UVOT UV and ASAS-SN $V$-band photometry (including detection limits) are plotted with open markers. The blue vertical lines on the bottom of the figure indicate the epochs of optical spectral observations, whereas the red ones indicate those of NIR spectral observations.  Table 1: Properties of ASASSN-15hy.aaAll photometric parameters are in the CSP-II natural system. ASASSN-15hy |
---|---
A.k.a. |  PS15aou
$\alpha$ (J2000) |  $20^{h}10^{m}02^{s}.35$
$\delta$ (J2000) |  $-00^{\circ}44^{\prime}21^{\prime\prime}.31$
$JD_{\text{max}}(B)$ |  2457151.6 $\pm$ 0.4
$JD_{\text{max}}(V)$ |  2457155.6 $\pm$ 0.7
$JD_{\text{max}}(r)$ |  2457157.1 $\pm$ 0.5
$JD_{\text{max}}(i)$ |  2457158.9 $\pm$ 1.1
$\Delta m_{15}(B)$ |  0.72 $\pm$ 0.04 mag
$\Delta m_{15}(V)$ |  0.44 $\pm$ 0.03 mag
$\Delta m_{15}(r)$ |  0.28 $\pm$ 0.02 mag
$\Delta m_{15}(i)$ |  0.12 $\pm$ 0.02 mag
$m_{B,\text{max}}$bbPeak apparent magnitude without any corrections. |  15.72 $\pm$ 0.01 mag
$m_{V,\text{max}}$bbPeak apparent magnitude without any corrections. |  15.38 $\pm$ 0.01 mag
$m_{r,\text{max}}$bbPeak apparent magnitude without any corrections. |  15.30 $\pm$ 0.01 mag
$m_{i,\text{max}}$bbPeak apparent magnitude without any corrections. |  15.50 $\pm$ 0.01 mag
$M_{B,\text{max}}$ccPeak absolute magnitude with $K$-correction, Milky Way and host-galaxy extinction correction. |  $-19.14^{+0.11}_{-0.16}$ mag
$M_{V,\text{max}}$ccPeak absolute magnitude with $K$-correction, Milky Way and host-galaxy extinction correction. |  $-19.36^{+0.11}_{-0.14}$ mag
$M_{r,\text{max}}$ccPeak absolute magnitude with $K$-correction, Milky Way and host-galaxy extinction correction. |  $-19.38^{+0.11}_{-0.13}$ mag
$M_{i,\text{max}}$ccPeak absolute magnitude with $K$-correction, Milky Way and host-galaxy extinction correction. |  $-19.10^{+0.11}_{-0.13}$ mag
$s_{BV}$ |  1.24 $\pm$ 0.18
$E(B-V)_{\text{MW}}$ |  0.13 mag
Host |
Heliocentric redshift |  0.0185 $\pm$ 0.0003
CMB redshift |  0.0176 $\pm$ 0.0003
Distance modulusddObtained from the CMB redshift, including the uncertainty due to peculiar velocity dispersion (assuming 300 km s-1). The cosmological parameters assumed for this work are $H_{0}=73$ km s-1 Mpc -1, $\Omega_{m}=0.27$ and $\Omega_{\Lambda}=0.73$.  |  34.33 $\pm$ 0.11 mag
$E(B-V)_{\text{host}}$ |  $0.00^{+0.03}$ mag
Stellar mass M⋆ |  (0.77 $\pm$ 0.16) $\times$ 109 M⊙
log10(sSFR) |  $-$10.21 $\pm$ 0.31 yr-1

ASASSN-15hy (aka PS15aou) was discovered by the All Sky Automated Survey for
SuperNovae (ASAS-SN, Shappee et al. 2014; Kochanek et al. 2017) in an image
taken with the double “Cassius” telescope on 2015 April 25.21 UT at $V$
$\sim$16.4 mag and was confirmed two days later in another image taken with
the Las Cumbres Observatory 1-m robotic telescope in Sutherland, South Africa
(Holoien et al., 2015). It was spectroscopically classified by the Public ESO
Spectroscopic Survey for Transient Objects (PESSTO; Smartt et al., 2015) as a
young peculiar SN Ia with similarities to SN 2003fg and SN 2006gz on 2015
April 27.36 UT (Frohmaier et al., 2015). On the same day as the
classification, at $\sim-$11.5 d relative to the rest-frame $B$-band maximum,
the CSP-II started an optical and near-infrared (NIR) follow-up campaign using
the 1-m Swope, 2.5-m du Pont, and the 6.5-m Baade telescopes, located at the
Las Campanas Observatory (LCO) in Chile. In this paper, we present
complementary multiwavelength photometric and spectroscopic observations of
ASASSN-15hy from the CSP-II, ASAS-SN, $Swift$, and Gemini. ASASSN-15hy is only
second 03fg-like object with a NIR spectroscopic time series and has a
multiband light-curve data set from UV to NIR. Note that the photometric
observation gap between +31 d and +96 d are due to the lack of available
observation time of the CSP-II during the Chilean winter. The photometry and
observation logs are listed in Appendix A.

###  2.1 Photometry

The ASAS-SN observations of ASASSN-15hy consist of $V$-band photometry,
including 3$\sigma$ nondetection limits. The ASAS-SN images were reduced with
an automated pipeline that uses the ISIS package (Alard & Lupton, 1998; Alard,
2000) for image subtraction, the IRAF222IRAF was distributed by the National
Optical Astronomy Observatory, which was managed by the Association of
Universities for Research in Astronomy (AURA) under a cooperative agreement
with the National Science Foundation. $apphot$ package for performing aperture
photometry, and the AAVSO Photometric All-sky Survey (APASS, Henden et al.
2015) for calibration.

CSP-II optical photometric observations in the $uBgVri$ bands of ASASSN-15hy
were obtained with the e2v 4112 $\times$ 4096 pixel CCD imager attached to the
Swope 1-m telescope at LCO. The NIR $YJH$ images were taken using RetroCam on
the LCO 2.5-m du Pont telescope. RetroCam employs a single chip Rockwell 1024
$\times$ 1024 HAWAII-1 HgCdTe detector.

The CSP-II photometric observations were reduced following the procedure
described in Phillips et al. (2019). Both optical and NIR images of
ASASSN-15hy were background subtracted, with optical and NIR host-galaxy
template images obtained at the Swope and Baade telescopes respectively, more
than +300 d past maximum light. Point-spread-function photometry was used to
compute the SN magnitude with respect to local-sequence stars, which were
calibrated with standard star fields. An $i$-band finder chart, with labeled
local-sequence stars, is shown in Fig. A16.

Finally, UV $uvw2$, $uvm2$, $uvw1$, and optical $ubv$ photometry was obtained
(PI: Holoien) with the Ultraviolet Optical Telescope (UVOT) onboard the
$Swift$ spacecraft (Gehrels et al., 2004; Roming et al., 2005). The
photometric reduction follows the same basic manner outlined by Brown et al.
(2009, 2014b).

Fig. 1 presents the observed light curves of ASASSN-15hy. All photometry
mentioned above is tabulated in Appendix A. The basic observational parameters
of ASASSN-15hy are summarized in Table 1. As much as possible, photometry in
the $BVri$ bands are placed via $S$-correction in the most recent CSP-II
natural system (Phillips et al., 2019), referred to as the “CSP natural
system” in the subsequent text.

###  2.2 Spectroscopy

Spectroscopic observations of ASASSN-15hy were obtained using multiple
telescopes and instruments. There are 34 optical spectra (including 22
previously published) from $-$11.6 d to +155.2 d and 6 NIR spectra from $-$8.5
d to +71.7 d relative to rest-frame $B$-band maximum, respectively. A journal
of the spectroscopic observations and the details of the published spectrum
are given in Appendix A.

Figure 2: Optical spectroscopic time series of ASASSN-15hy of a selected
sample. Regions of strong telluric absorptions are marked with vertical gray
bands. The colors represent various telescopes and instruments. Note that some
edges of the spectra are grayed out for presentation purposes due to the high
level of noise.  Figure 3: NIR spectroscopic time series of ASASSN-15hy. The
strongest telluric absorptions are marked with vertical gray bands.

This work contribute twelve optical spectroscopic observations of ASASSN-15hy.
Five of them were obtained with the FAST spectrograph (Fabricant et al., 1998)
on the 1.5-m Tillinghast telescope at the F. L. Whipple Observatory (FLWO 1.5
m) and were reduced using a combination of custom IDL and standard IRAF
procedures (Matheson et al., 2005). Four of the optical spectra were taken
with the SPRAT spectrograph on the 2-m Liverpool Telescope (LT, Steele et al.
2004) and reduced using the LT pipeline (Piascik et al., 2014). Two spectra
obtained with ALFOSC on the Nordic Optical 2.5-m Telescope (NOT) and one
spectrum taken with WFCCD mounted on the 2.5-m du Pont telescope were all
reduced in the standard manner using IRAF scripts. The optical spectroscopic
observations of ASASSN-15hy are shown in Fig. 2.

The NIR spectral observations were obtained using GNIRS (Elias et al., 1998)
on the 8.2-m Gemini North Telescope (GN), FLAMINGOS-2 (F2, Eikenberry et al.
2008) on the 8.2-m Gemini South Telescope (GS), and FIRE (Simcoe et al., 2013)
on the 6.5-m Baade Telescope. The observing techniques and reduction
procedures of these three instruments are described in Hsiao et al. (2019).
The NIR spectra of ASASSN-15hy are presented in Fig. 3.

###  2.3 Host Observations

Integral-field spectroscopy of the host galaxy of ASASSN-15hy was obtained
when the SN had faded, on 2017 July 19 UT, with the Multi-Unit Spectroscopic
Explorer (MUSE, Bacon et al. 2010), mounted on the Unit 4 telescope at the ESO
Very Large Telescope (VLT UT4) of the Cerro Paranal Observatory. These host
observations were obtained as part of the All-weather MUSE Supernova Integral-
field Nearby Galaxies (AMUSING, Galbany et al. 2016a) survey, an ongoing
project studying the environments of SNe by means of the analysis of a large
number of nearby SN host galaxies.

MUSE provides a wide field-of-view of approximately 1${}^{\prime}\times$1′ and
a square spatial pixel size of 0.2″$\times$0.2″. This limits the spatial
resolution of the data to the atmospheric seeing during the observations,
which was 1.6″ while observing the host galaxy of ASASSN-15hy. The
spectroscopy covers a wavelength range from 4750 to 9300 Å, with a mean
spectral resolution $\lambda$/$\Delta\lambda$ $\sim$3000 and a spectral
sampling of 1.25 Å.

##  3 Host Properties

The host-galaxy properties of 03fg-like SNe can be extreme, and as a group,
distinct from those of normal SNe Ia. Detailed studies of these galaxies are
therefore likely pertinent to understanding the origin of 03fg-like SNe. With
the integral-field spectroscopy taken by MUSE, we conducted a similar analysis
to that presented by Hsiao et al. (2020) in order to obtain several
measurements of galaxy properties, both at the location of ASASSN-15hy and for
the galaxy as a whole.

Figure 4: Host-galaxy images of ASASSN-15hy extracted from the MUSE datacube.
The left panel corresponds to the synthetic CSP $r$-band image, and the right
panel is the extinction-corrected and stellar population subtracted gas-phase
emission line map of H$\alpha$. The SN position is marked with a red cross in
both panels, and the host galaxy is defined by a green contour (the bright
point source inside the contour in the left panel is a foreground star). The
numbers from 1 to 6 mark the field stars that show residuals in the H$\alpha$
map. Letter A in the right panel marks the position of the host structure
where analysis is performed (see Appendix C.1).

###  3.1 Host Redshift

In most Swope images of ASASSN-15hy, the SN appears hostless, which is
remarkable given that the estimated redshift based on the SN spectrum is
$z=0.025$, which is relatively nearby (Frohmaier et al., 2015). With the
deeper MUSE data, a faint and extended source is apparent in the synthetic
$r$-band image (left panel of Fig. 4), produced by convolving the transmission
function of the CSP $r$-band filter with the MUSE datacube. The extended
source observed to the south east (SE) of the SN is presumed to be the host of
ASASSN-15hy. We manually defined a contour (marked with green in Fig. 4) that
includes the extended source but excludes a foreground star. Note that another
bright point source within the contour in the synthetic $r$-band image is a
foreground star and is excluded in the analysis. The extracted spectrum within
this contour revealed the typical emission lines of a star-forming galaxy.

A heliocentric redshift for the host galaxy of $z_{\text{helio}}$ = 0.0185
$\pm$ 0.0003 was derived by measuring the positions of the strongest emission
lines, such as H$\alpha$, [O iii] 5007 Å, and H$\beta$. Correcting to the
cosmic microwave background (CMB) rest frame yielded $z_{\text{CMB}}$ = 0.0176
$\pm$ 0.0003, which translates to a distance modulus ($\mu$) of 34.33 $\pm$
0.11 mag assuming $H_{0}=73$ km s-1 Mpc -1, $\Omega_{m}=0.27$ and
$\Omega_{\Lambda}=0.73$ for the cosmological parameters. Note that the quoted
uncertainty in $\mu$ includes an assumed peculiar velocity dispersion of 300
km s-1.

###  3.2 Host Environment

To study the host properties of ASASSN-15hy, the MUSE data were analyzed using
methods similar to those of Galbany et al. (2014, 2016b) and Hsiao et al.
(2020). The right panel of Fig. 4 shows the extinction-corrected H$\alpha$
emission map of the host and reveals structures across the galaxy that are not
seen in the synthetic $r$-band image shown in the left panel of Fig. 4.
Spectra were extracted at several locations, and simple stellar population
analysis was performed in order to obtain several measurements characterizing
the host-galaxy properties. These measurements are listed in Table 2.

The host of ASASSN-15hy has a low stellar mass and extremely low metallicity
indicated by the subsolar oxygen abundance. It contains a remarkably young
stellar population component based on the high H$\alpha$ equivalent width.
Furthermore, this low-mass galaxy is very efficient in producing new stars, as
shown by the high specific star formation rate (sSFR). The details of the host
environment analysis and the results are presented in Appendix C.1.

In the context of 03fg-like SNe, the properties of the host environment of
ASASSN-15hy are largely consistent with other group members. The MUSE
observation of the LSQ14fmg host shows similar properties, such as low
metallicity, low stellar mass, high sSFR, and a relatively young stellar
population component (Hsiao et al., 2020). The host of SN 2007if was also
shown to be a low-stellar-mass and metal-poor galaxy (Scalzo et al., 2010;
Childress et al., 2011). The host environment of 03fg-like objects indicates
that they are likely to originate from low-metallicity and young progenitors
(Childress et al., 2011; Khan et al., 2011; Silverman et al., 2011;
Taubenberger et al., 2011). However, the current small sample and the lack of
a complete and consistent theoretical scenario hinder our ability to
disentangle these effects and identify the dominant environmental driver for
this class of objects.

Table 2: Properties of ASASSN-15hy and the host environment from the MUSE data. Location | H$\alpha$EW | 12 + $log_{10}(O/H)$ | Stellar mass | SFR | log10(sSFR)
---|---|---|---|---|---
| (Å) | O3N2 (dex) | D16 (dex) | (log10(M⋆/M⊙)) | (M⊙ yr-1) | (yr-1)
SN | 30.48$\pm$0.89 | 8.43 $\pm$ 0.11 | 8.04 $\pm$ 0.35 | 6.77 $\pm$ 0.10 | 0.20$\pm$0.09$\times$10-3 |  $-$10.47 $\pm$ 0.61
Structure A | 75.30$\pm$0.48 | 8.21$\pm$0.14 | 7.84$\pm$0.15 | $\dots$ | 1.21$\pm$0.28$\times$10-3 | $\dots$
Global | 33.25$\pm$0.28 | 8.34$\pm$0.14 | 7.87$\pm$0.16 | 8.89 $\pm$ 0.09 | 0.048$\pm$0.007 |  $-$10.21 $\pm$ 0.31

###  3.3 Host-galaxy Extinction

Obtaining an accurate value of the host-galaxy extinction of 03fg-like SNe is
difficult. As described by Scalzo et al. (2012), the Lira relation (Lira,
1996; Phillips et al., 1999) does not hold for these objects. Nevertheless,
Chen et al. (2019) derived a color excess for SN 2009dc by utilizing the Lira
method and assuming that ASASSN-15pz has zero host-galaxy reddening. This
strategy relies on the assumption that the difference in the late-time $B-V$
color is not intrinsic to the SN Ia. In this work, the relation between
$E(B-V)$ and the Na i D equivalent width (EW) from Poznanski et al. (2012) and
Phillips et al. (2013) was adopted to estimate the host reddening and
extinction of ASASSN-15hy.

Each optical spectrum of ASASSN-15hy was inspected for a narrow Na i D
absorption at the redshift of the host galaxy, and no absorption features were
found. A detection limit of EW $\leq$ 0.1 Å was estimated using the WiFeS
spectrum taken on 2015 May 13, selected for its relatively high resolution and
signal-to-noise ratio (S/N). The detail of the methodology is described in
Appendix C.2. The Na i D absorption from the Milky Way (MW) was measured to
have an EW of 0.82 $\pm$ 0.25 Å in the same spectrum, corresponding to a MW
color excess of $E(B-V)_{\text{MW}}=0.13\pm 0.09$ mag based on the relation
presented in Poznanski et al. (2012). This is consistent with the value
obtained from the Galactic reddening map of Schlafly & Finkbeiner (2011).

From the nondetection of the host-galaxy Na i D lines, we assume zero
extinction throughout the paper. The extinction estimate from the Balmer
decrement of the MUSE observation at the location of the SN is also consistent
with zero. An extinction uncertainty $\delta A_{V}$ = $0.08$ mag is estimated
based on the MW Na i D and extinction measurements compiled by Phillips et al.
(2013) and references therein, using the Na i D EW detection limit of 0.1 Å.
This translates to an uncertainty of 0.03 mag for the color excess. Hence,
throughout this work, a host-galaxy color excess of
$E(B-V)_{\text{host}}=0.00^{+0.03}$ mag is adopted for ASASSN-15hy.

ASASSN-15hy is dim in the optical and red in $B-V$ near maximum in comparison
to the 03fg-like group (see Section 4) and begs the question whether these
properties are intrinsic to the SN or due to host-galaxy reddening.
Comparisons of the spectral energy distributions (SEDs) of ASASSN-15hy with
those of well-observed 03fg-like SNe Ia show that these properties are likely
to be intrinsic to the SN. To match the optical flux of ASASSN-15hy to those
of ASASSN-15pz or SN 2009dc by invoking host-galaxy extinction requires an
enormous intrinsic UV flux from ASASSN-15hy that is $2-5$ times higher than
the UV bright SN 2009dc and ASASSN-15pz. The nondetection of Na i D reaffirms
that the SED differences are dominated by differences in the intrinsic
properties. Details of the SED comparisons are presented in Appendix C.3.

##  4 Photometric Properties

The observed light curves of ASASSN-15hy are presented in Fig. 1. In this
section, the photometric properties of ASASSN-15hy are analyzed by comparing
light-curve parameters, such as the $B$-band decline
rate333$\Delta{\rm{m}_{15}}(B)$ is the change in magnitude between maximum
light and +15 d past maximum in a certain band (Phillips, 1993).
$\Delta{\rm{m}_{15}}(B)$ and the color stretch444$s_{BV}$ is the difference in
time between $B$-band maximum and the reddest point in the $B-V$ color curve,
normalized by 30 days (Burns et al., 2014). $s_{BV}$, as well as the shapes of
the light and color curves relative to other SNe Ia. The basic observational
parameters of ASASSN-15hy are summarized in Table 1.

Six of seven published 03fg-like objects so far: SN 2006gz (Hicken et al.,
2007), SN 2007if (Scalzo et al., 2010; Friedman et al., 2015; Krisciunas et
al., 2017), SN 2009dc (Taubenberger et al., 2011; Hicken et al., 2012;
Friedman et al., 2015; Krisciunas et al., 2017), SN 2012dn (Chakradhari et
al., 2014; Yamanaka et al., 2016; Taubenberger et al., 2019), ASASSN-15pz
(Chen et al., 2019), and LSQ14fmg (Hsiao et al., 2020), are included for
comparison in this section. Note that SN 2003fg (Howell et al., 2006) is not
included in the comparison due to the lack of $B$\- and $V$-band light curves.
A well observed normal SN Ia, 2007af (Krisciunas et al., 2017), is included
for comparison in this section. All photometry is $K$-corrected (and
S-corrected if the filter function is available) following the manner
described in Appendix B.

These SNe were selected for comparison to ASASSN-15hy as they allow for the
full diversity of properties to be explored, and other subclasses (such as
2002cx-like, 1991bg-like) are not included here due to highly distinct
differences with 03fg-like objects (see Fig. 1 of Taubenberger 2017).

Figure 5: The luminosity width relation (Phillips relation) for normal SNe Ia
from the CSP-I (Krisciunas et al., 2017) is demonstrated by gray markers.
Published 03fg-like objects (see references in text) are plotted in colored
markers. All absolute peak magnitudes are corrected for MW extinction and
host-galaxy extinction.

###  4.1 Luminosity Width Relation

SNe Ia follow a luminosity width relation (LWR), where SNe Ia with broader
light curves are intrinsically brighter, and SNe with narrower light curves
are dimmer (Pskovskii, 1977; Phillips, 1993; Phillips et al., 1999). This
relation serves as a cornerstone of modern SN Ia cosmology. To determine where
ASASSN-15hy and other 03fg-like SNe are located on the luminosity-width
diagram, the $K$-corrected $B$-band light curves were used to directly measure
the time of maximum, decline rate $\Delta{\rm{m}_{15}}$(B), and the apparent
peak magnitude $m_{B,\rm{max}}$. These measurements were obtained through a
Gaussian process interpolation included in the SuperNovae in object-oriented
Python (SNooPy; Burns et al. 2011, 2014) package. Finally, MW and host-galaxy
extinction corrections and distance modulus were applied to the
$m_{B,\rm{max}}$ to obtain the absolute peak magnitude $M_{B,\rm{max}}$,
hereafter $M_{B}$ for short.

ASASSN-15hy reached $B$-band maximum at $JD_{\rm{max}}(B)$ = 2457151.6 $\pm$
0.4, has a $\Delta{\rm{m}_{15}}(B)$ of 0.72 $\pm$ 0.04 mag, and a
$m_{B,\rm{max}}$ of 15.72 $\pm$ 0.01 mag. As suggested in Section 3.3, there
is no evidence of ASASSN-15hy suffering from substantial host-galaxy
extinction. Hence, only a MW color excess $E(B-V)_{\text{MW}}=0.13$ mag was
used to correct the photometry of ASASSN-15hy, which translates to a $B$-band
extinction of AB = 0.53 mag based on the extinction law from Cardelli et al.
(1989) assuming RV = 3.1. Converting to absolute magnitude, using the AB value
above, and a $\mu$ of 34.33 $\pm$ 0.11 mag from the host-galaxy redshift in
the CMB frame gives $M_{B}=$ $-$19.14 $\pm$ 0.11 mag. This corresponds to a
$M_{B}$ of $-19.14^{+0.11}_{-0.16}$ mag if including the uncertainty of the
host extinction.

Table 3: Distance modulus and color excess values adopted for comparison
sample of 03fg-like SNe Ia.

SN | $z_{\text{CMB}}$ |  $\mu$aaObtained from the CMB redshift with the same cosmological parameters assumed for ASASSN-15hy (see Section 3.1 ). |  $E(B-V)$ccObtained from Schlafly & Finkbeiner (2011). |  $E(B-V)$ddAdopted from Hicken et al. (2007), Scalzo et al. (2010), Taubenberger et al. (2011), Taubenberger et al. (2019), Chen et al. (2019), Hsiao et al. (2020) and this work respectively.
---|---|---|---|---
|  | (mag) | (MW, mag) | (host, mag)
2006gz | 0.023 | 34.95$\pm$0.09 | 0.05 | 0.18$\pm$0.05
2007if | 0.073 | 37.51$\pm$0.03 | 0.07 | 0.0
2009dc | 0.022 | 34.79$\pm$0.09 | 0.06 | 0.10$\pm$0.07
2012dn | 0.009 | 33.28$\pm$0.21bbCorrected for infall towards the Virgo cluster and the Great Attractor (recession velocity = $3{,}306$ km s-1, Mould et al. 2000). | 0.05 | 0.04$\pm$0.01
ASASSN-15pz | 0.014 | 33.89$\pm$0.14 | 0.01 | 0.0
LSQ14fmg | 0.065 | 37.24$\pm$0.03 | 0.05 | 0.13$\pm$0.15
ASASSN-15hy | 0.0176 | 34.33$\pm$0.11 | 0.13 | $0.00\pm 0.03$

Figure 5 demonstrates where ASASSN-15hy and other 03fg-like events are located
on the luminosity-width diagram. ASASSN-15hy is intrinsically the least
luminous 03fg-like SN and is located below the LWR for normal SNe Ia; whereas
SN 2006gz, SN 2012dn, and ASASSN-15pz blend into the normal population; and SN
2007if, SN 2009dc, and LSQ14fmg are all overluminous compared to normal SNe
Ia. The corresponding adopted $\mu$ and reddening values of the 03fg-like SNe
are tabulated in Table 3. A $\mu$ of 32.09 $\pm$ 0.32 mag and
$E(B-V)_{\text{host}}=0.19~{}\pm~{}0.01$ mag (obtained from $SNooPy$) are
adopted for the normal SN Ia 2007af.

ASASSN-15hy is almost one magnitude dimmer than SN 2009dc despite the fact
that they have very similar $\Delta{\rm{m}_{15}}(B)$ values, NIR luminosities
(see Section 4.2), and spectral properties (see Section 5). If ASASSN-15hy
followed the LWR, for its absolute $M_{B}$ magnitude of $\sim-$19.14 mag, it
would be expected to have a $\Delta{\rm{m}_{15}}(B)$ of $\sim 1.4$ mag,
demonstrating the uniqueness of this event. Cosmology light-curve fitters are
liable to give incorrect results for SNe similar to ASASSN-15hy because of the
red intrinsic color at maximum (see Section 4.3).

(a) Time axis relative to peak time in each band. (b) Time axis relative to
peak time in $B$ band.

Figure 6:  Comparison of $K$-corrected $BVri$ light curves. The magnitudes are
normalized to the peak magnitude in each band, and the rest-frame time axes
are presented relative to the peak time of each respective band in (a) and to
the peak time in the $B$ band in (b). The black lines are the Ia template
light-curves fits ASASSN-15hy using SNooPy, treating it as if it were a normal
SN Ia.

###  4.2 Light-curve Evolution

Figure 6 compares the optical rest-frame $K$-corrected $BVri$ light curves of
ASASSN-15hy and other SNe in the comparison sample described earlier. A
91T-like SN Ia, LSQ12gdj (Scalzo et al., 2014) is added to the sample for
comparison. Another peculiar object, SN 2006bt (Foley et al., 2010;
Stritzinger et al., 2011; Krisciunas et al., 2017), which is underluminous but
slowly declining, is also included in the comparisons since it has a similar
$i$-band light-curve shape to 03fg-like SNe Ia. The $B$\- and $V$-band light
curves of all SNe in the figure have similar shapes but differ in time scale.
ASASSN-15hy evolves more slowly on the rise compared to the SNooPy SN Ia
template, but behaves similarly to some other 03fg-like SNe. The rise time of
ASASSN-15hy (22.5 $\pm$ 4.6 d), determined by fitting the flux converted from
the CSP-II and ASAS-SN $V$-band pre-maximum magnitudes with a second order
polynomial, is similar to most of the 03fg-like objects where values can be
determined. In this respect, ASASSN-15hy is particularly similar to SN 2007if
(Scalzo et al., 2010) and SN 2009dc (Silverman et al., 2011; Taubenberger et
al., 2011), which have rise times of 24.2 $\pm$ 0.4 d and 23 $\pm$ 2 d, and
decline rates $\Delta{\rm{m}_{15}}(B)$ of 0.71 $\pm$ 0.03 mag and 0.71 $\pm$
0.06 mag respectively. From maximum light to $\sim$+30 d, ASASSN-15hy has
similar behavior to the comparison SNe. However, after 30 d past maximum
light, ASASSN-15hy has among the slowest declining light curves in all $BVri$
bands. In the $B$ band, ASASSN-15hy declines by $\sim$1.0 mag from +30 to +90
d, whereas SN 2009dc declines by $\sim$1.4 mag, for example.

The $r$\- and $i$-band light curves of ASASSN-15hy are similar to other 03fg-
like objects but differ significantly from other SNe Ia. In the $r$ band,
normal and 91T-like SNe show a shoulder feature at $\sim$+15 d after peak,
while ASASSN-15hy and other 03fg-like objects generally do not show this
feature. In the $i$ band, ASASSN-15hy has a slower rise compared to normal and
91T-like SNe Ia and lacks a pronounced secondary maximum. Note that SN 2006bt
also lacks a clear secondary maximum in the $i$ band, however, it differs in
NIR light-curve properties and spectra compared to 03fg-like objects (see
below). ASASSN-15hy has an $i$-band primary maximum that occurs after the
$B$-band maximum and has a difference in peak times ($t^{\rm{peak}}_{i-B}$) of
7.2 $\pm$ 1.1 d. Although this is a ubiquitous property of 03fg-like objects
(Ashall et al., 2020), ASASSN-15hy has the largest value of
$t^{\rm{peak}}_{i-B}$ among all 03fg-like objects included in Ashall et al.
(2020), which have a mean of 3.0 $\pm$ 2.1 d.

In fact, the overall trend of $BVri$ peak times and decline rates indicates
the extreme nature of ASASSN-15hy. ASASSN-15hy declines progressively more
slowly and peaks progressively later in time, from blue to redder bands. This
trend is significantly different from the normal SN Ia 2007af, which does not
show the same progressive change especially from $r$ to $i$ band. On the other
hand, the 03fg-like objects all show a relatively smooth transition, similar
to ASASSN-15hy. However, ASASSN-15hy is on the extreme end in its light-curve
properties, including the slowest decline rate and the largest difference in
peak times in both $r$ and $i$ bands.

To compare the intrinsic brightnesses of ASASSN-15hy and others, Fig. 7
presents the absolute magnitudes, corrected for MW and host extinction, of the
comparison SNe Ia in the $uvm2$ and $uBYJH$ bands. The larger error bars are
due to the high uncertainties in the host reddenings for some of the
comparison SNe Ia. UVOT photometry of the comparison SNe was obtained from the
$Swift$ Optical Ultraviolet Supernova Archive (SOUSA; Brown et al. 2014a) via
the Swift Supernova website.555https://pbrown801.github.io/SOUSA/

ASASSN-15hy is UV bright and has a peak absolute magnitude of
$-18.17^{+0.11}_{-0.27}$ mag in the $uvm2$ band, which is more than 1 mag
brighter comparing to the normal SN Ia 2007af around the same time, consistent
with other 03fg-like objects (Taubenberger et al., 2011; Brown et al., 2014b;
Chen et al., 2019). The $uvm2$ light curve of ASASSN-15hy is most comparable
to SN 2012dn, and they both peak earlier than the normal SN Ia 2007af. The
early bright UV feature of 03fg-like SNe may be explained by the shock
interaction (e.g., Fryer et al., 2010; Blinnikov & Sorokina, 2010;
Taubenberger et al., 2011; Hachinger et al., 2012; Scalzo et al., 2012).

During the photospheric phase, ASASSN-15hy is one of the faintest 03fg-like
SNe in the $u$ and $B$ bands at $B$-band maximum. It is $\sim$1.4 and
$\sim$1.0 mag fainter than SN 2009dc in the $u$ and $B$ band, respectively. In
fact, ASASSN-15hy is fainter than the normal SN Ia 2007af, which has a faster
decline rate. Interestingly, ASASSN-15hy has a similar brightness to other
03fg-like SNe in the NIR bands.

As shown in Fig. 7, the NIR light curves of ASASSN-15hy are broader, brighter
than normal SNe Ia and do not have two distinct maxima, similar to other 03fg-
like objects. In the $Y$ and $H$ bands, ASASSN-15hy continued to rise in
brightness over the duration of our NIR observations and does not appear to
have reached peak by the final observations at $\sim$+22 d past $B$-band
maximum, which have magnitudes of $-19.45$ and $-$19.39 mag respectively. The
03fg-like events SN 2009dc, SN 2012dn and ASASSN-15pz all show weak secondary
maxima around +30 d in the $J$ band, similar to the timing of the NIR
secondary maximum for a normal SN Ia. Unfortunately, our $J$-band light curve
does not extend to these epochs.

SN 2006bt shares some similar properties in the optical with 03fg-like
objects, but shows a different evolution in the NIR; for example, SN 2006bt
showed a decrease in flux in the $Y$ band after the $B$-band peak. Clearly,
NIR light curves can be used as a powerful diagnostic to distinguish 03fg-like
SNe from other subclasses of SNe Ia.

Figure 7: Comparison of the absolute magnitudes in UV $uvm2$, optical $uB$ and
NIR $YJH$. The absolute magnitudes are computed from $M=m-\mu-
A_{\text{MW}}-A_{\text{host}}$. Except for the $B$ band, light curves in this
figure are not $K$-corrected due to the lack of spectroscopic coverage.

###  4.3 Color Curves

Figure 8: Comparison of color curves. $B-V$, $r-i$, and $g-r$ color curves
were obtained from $K$-corrected $BgVri$ photometry. The color curves in the
UV and NIR bands are not $K$-corrected due to the lack of spectroscopic
coverage. All data presented are MW extinction corrected.

The observed $uvm2-u$, $B-V$, $r-i$, $g-r$, and $J-H$ color curves of
ASASSN-15hy are presented in Fig. 8, along with those of other 03fg-like SNe
and the normal SN Ia 2007af. Here, host extinction corrections, which are
themselves uncertain, are not applied to all comparison SNe.

In $B-V$ and $g-r$, ASASSN-15hy shows the reddest observed colors of all 03fg-
like SNe at maximum light and has values of $0.18\pm 0.01$ mag and $0.07\pm
0.01$ mag respectively, compared to an range of $-$0.04 to 0.09 mag and
$-$0.21 to $-$0.07 mag for other 03fg-like SNe. This color difference would be
even more pronounced after correcting for host-galaxy reddening, since
ASASSN-15hy has negligible host extinction. These colors of ASASSN-15hy evolve
monotonically redward from the first observation at $-11.5$ d relative to the
$B$-band maximum until $\sim$+30 d. While most of other 03fg-like SNe show
local color minima around 0 d similar to the normal SN Ia 2007af, except for
SN 2007if (it is unclear due to the lack of observations). If ASASSN-15hy has
such local minimum, the timing of the bluest point would be much earlier than
most of the other 03fg-like SNe. Although any possible host-galaxy extinction
may shift the $B-V$ color curve vertically, the shape of the color curve would
not change significantly. Thus, ASASSN-15hy is unique in that it does not have
an initial phase where it evolves bluer in $B-V$ and $g-r$, as well as the
reddest color among 03fg-like SNe at $B$-band maximum. Our observations may
not be early enough to capture the bluest point of $B-V$ and $g-r$;
nevertheless, this turn over point is shown to be possibly the earliest among
03fg-like SNe.

The color stretch parameter, $s_{BV}$, can also be used to characterize a SN
Ia and is determined by locating the timing of the reddest point in the $B-V$
color curve (Burns et al., 2014). ASASSN-15hy has an $s_{BV}$ of 1.24 $\pm$
0.18 (the large error due to the lack of photometric coverage at the reddest
point). ASASSN-15hy, along with all other 03fg-like SNe, have $s_{BV}$ values
larger than 1.0, matching the slowest decliners in the normal population. The
large $s_{BV}$ together with a late $t^{\rm{peak}}_{i-B}$ are uniform features
of these 03fg-like SNe, enabling classification between peculiar subgroups of
SNe Ia (Ashall et al., 2020).

The $r-i$ color curve of ASASSN-15hy is also unique compared to normal SNe Ia
and the majority of 03fg-like events. As Hsiao et al. (2020) have shown, the
$r-i$ color curve is one of the more obvious ways to distinguish this subclass
from normal objects. Compared with normal SNe Ia, these peculiar 03fg-like
objects evolve more slowly in time, as well as having shallower peaks and
troughs in their $r-i$ evolution. These differences are most likely caused by
the unusually broad $i$-band light curve in 03fg-like objects that also
display very weak or missing $i$-band secondary maxima. Within the 03fg-like
group, there is a large scatter in $r-i$ colors at the epoch of $B$ maximum.
ASASSN-15hy is separated from the majority of 03fg-like members and is at
least $\sim$0.1 mag redder than other 03fg-like objects around the time of
$r-i$ minimum, but is not as extreme as LSQ14fmg. The unusually flat $r-i$
color of LSQ14fmg is suggested to be the outcome of interaction of the SN
ejecta with a superwind of an AGB star (Hsiao et al., 2020), which could
potentially be the case for ASASSN-15hy as well.

The top panel of Fig. 8 shows the $Swift$ UVOT $uvm2-u$ color curves.
ASASSN-15hy evolves to the red up to $\sim$+20 d and then turns blueward. This
is similar to the other members of the 03fg-like group (e.g., SN 2009dc and
2012dn), and ASASSN-15hy appears to be the bluest at early times. The blue
UV$-$optical color ($uvm2-u$ = $1.72\pm 0.16$ mag at $B$-band maximum) may
indicate a lack of line-blanketing from Fe lines comparing to other members,
possibly caused by differences in metallicity (e.g., Lentz et al. 2000) or
caused by differences in the outer slope of the density profile of the ejecta
(Walker et al., 2012). Conversely, the $uvm2-u$ color of the normal SN Ia
2007af starts quite red and evolves monotonically blueward over time. 03fg-
like SNe are known to be UV-blue as well as UV-bright at early times (Milne et
al., 2013; Brown et al., 2014b; Chen et al., 2019), and this unique early UV
evolution separates them from the normal population.

The $J-H$ color curve of ASASSN-15hy is similar to that of SN 2009dc and other
03fg-like SNe but evolves differently when compared to the normal SN Ia
2007af. This is not a surprise given the difference between the NIR light
curves of 03fg-like and normal SNe Ia. For example, normal SNe Ia become
rapidly redder in $J-H$ until $\sim$+15 d past $B$-band maximum, then drop to
a local blue minimum at $\sim$+30 d. They then evolve to the red again as the
observations show. The $J-H$ red bump $\sim$+15 d is thought to be associated
with the evolution of the $H$-band break and the unveiling of iron group
elements in the ejecta (Wheeler et al., 1998). Unlike normal SNe Ia,
ASASSN-15hy becomes redder more slowly over time and only increases $\sim$0.3
mag from 0 to +20 d, whereas SN 2007af increases $\sim$1.5 mag during the same
time period. SN 2009dc and SN 2012dn also show a similar slow reddening at
early times and then flatten off in color between +15 to +40 d, possibly due
to the lack of a $H$-band break in the NIR spectra. It is noteworthy that SN
2007if evolves bluer in $J-H$ after +15 d until the last observed epoch,
although the uncertainties are large.

(a) Pre-maximum comparison to various subgroups. (b) Post-maximum comparison
within the 03fg-like subgroup.

Figure 9: Spectroscopic comparisons of ASASSN-15hy and other SNe Ia. Panel (a)
compares ASASSN-15hy with a normal SN Ia and other subgroups, including SN
2011fe (Zhang et al., 2016), SN 2009dc (Taubenberger et al., 2011), SN 1991T
(Filippenko et al., 1992a), SN 2006bt (Foley et al., 2010), SN 2012Z
(Stritzinger et al., 2015) and SN 2015bo (Yaron & Gal-Yam, 2012), around a
week before $B$-band maximum. The gray spectrum in the background is
ASASSN-15hy at $-9.3$ d, and the spectra in color are the comparison SNe Ia
around a similar phase. Panel (b) compares ASASSN-15hy with other published
03fg-like objects around a week after $B$-band maximum. The gray spectrum in
the background is ASASSN-15hy at $+4.5$ d, the purple ones are other 03fg-like
SNe Ia around a similar phase, and the normal SN Ia 2011fe is included for
comparison as well. The telluric features are marked using the color
corresponding to the SN.

##  5 Spectroscopic Properties

The spectroscopic data set of ASASSN-15hy has broad wavelength coverage and
high cadence. In total, there are 42 optical ($-$11.6 to $+$155.2 d) and 6 NIR
($-$8.5 to $+$71.7 d) spectroscopic observations, enabling a detailed look at
the spectral features and evolution of a rare 03fg-like object. Note that all
phases in this section are relative to the epoch of $B$-band maximum and
corrected for time dilation, unless otherwise stated.

###  5.1 Optical Wavelengths

The first two optical spectra were taken concurrently with the first Swope
multifilter observations (Fig. 2) and appear blue as confirmed by the $B-V$
color (Fig. 8). The SED then evolves redwards rapidly. The early optical
spectra are dominated by P-Cygni profiles of intermediate-mass elements (IME;
such as Si, Ca, and S), unburned elements (C and O), and iron group elements
(IGEs). Most of these features are also prevalent in the spectra of normal SNe
Ia. However, there is a lack of higher-ionization species such as Fe iii in
ASASSN-15hy, which are typical of normal, slow-declining SNe Ia.

####  5.1.1 Comparison of Spectral Features

The left panel of Fig. 9 presents a pre-maximum optical spectrum of
ASASSN-15hy with labeled line identifications, as well as comparisons between
ASASSN-15hy and a variety of subclasses of SNe Ia. The ions present in
ASASSN-15hy are similar to those observed in the normal SN 2011fe. The notable
differences are the overall weaker spectral features in ASASSN-15hy,
particularly the Ca ii features as well as the stronger C ii features. The
weak spectral lines of ASASSN-15hy resemble those of SN 1991T and SN 2012Z
(02cx-like). However, the ionization states of these two objects are much
higher than ASASSN-15hy, as the early spectra of SN 1991T and SN 2012Z are
dominated by Fe iii lines (Filippenko et al., 1992a; Stritzinger et al.,
2015). ASASSN-15hy does not exhibit the Ti ii feature around 4000 Å, which is
present in 06bt- and 91bg-like SNe (e.g., Filippenko et al. 1992b; Foley et
al. 2010). Overall, the spectra of ASASSN-15hy are remarkably similar to the
03fg-like event SN 2009dc (Taubenberger et al., 2011), demonstrating that this
object belongs in the 03fg-like subclass.

The right panel of Fig. 9 presents a comparison between ASASSN-15hy and other
03fg-like SNe around one week after $B$-band maximum. All of the objects have
a similar ionization state, are dominated by IMEs, show weak Ca ii H&K
features, and most have persistent C ii absorptions past maximum light.
However, the strength and velocity of the Si ii $\lambda$6355 feature vary
within this subclass.

SN 2009dc and ASASSN-15hy are the two 03fg-like SNe Ia that have the most
complete optical spectroscopic data sets. Figure 10 shows the time series
comparison between these two 03fg-like objects. Again, these two objects have
very similar evolution in their spectral features, except for the slightly
redder color of ASASSN-15hy. The C ii $\lambda$6850 line is still present in
both ASASSN-15hy and SN 2009dc around one week after maximum, while C ii in a
normal SN Ia, if present at all, usually disappears soon after the explosion
(Taubenberger et al., 2011; Folatelli et al., 2012). The strong and persistent
C ii feature implies substantial unburned material in the ejecta (Howell et
al., 2006; Hicken et al., 2007).

Figure 10: Spectroscopic time series comparison of ASASSN-15hy and SN 2009dc
in the optical. Spectra of ASASSN-15hy are plotted with purple colors and
those of SN 2009dc (Taubenberger et al., 2011) around similar phases are
plotted with gray in the background. Regions of strong telluric adsorptions in
the rest frame of ASASSN-15hy are marked with the purple symbols at the top.

####  5.1.2 Si ii and C ii Evolution

To compare the line profiles of ASASSN-15hy and other SNe quantitatively, we
measured the velocities and pseudo-equivalent widths (pEW) of the Si ii
$\lambda$6355 and C ii $\lambda$6850 features. We first divide the flux by the
local continuum around the feature, which is constructed by connecting the two
local maxima on both sides of the feature. The area is integrated to obtain
the pEW. The position of the absorption minimum is obtained by fitting a
Gaussian profile to obtain the blue-shifted velocity. To estimate the
measurement uncertainties, we bootstrap-resample the flux 500 times repeating
the above process. The mean and standard deviation of the realizations are
adopted as the measured value and corresponding error for each pEW and
velocity measurement.

ASASSN-15hy has small pEW values for Si ii $\lambda$5972 and $\lambda$6355
near maximum light, which places it in the “shallow silicon” (SS) group on the
Branch diagram (Branch et al., 2006). This is a common characteristic of 03fg-
like events (e.g., Taubenberger et al. 2011; Hsiao et al. 2020), as shown in
Fig. 11. The only 03fg-like SN that might sit in a different area of the
Branch diagram is SN 2012dn; however, this SN also has a lower luminosity than
most of the other 03fg-like SNe. Interestingly, the underluminous ASASSN-15hy
is sitting closer to the brighter 03fg-like SNe, such as SNe 2007if and
2009dc, on the Branch diagram.

The velocity of Si ii $\lambda$6355 for ASASSN-15hy ranges from $\sim 9{,}300$
km s-1 at -10 d to $\sim 8{,}600$ km s-1 at +2 d and is $\sim 2{,}000$ km s-1
slower than the normal SN 2011fe. The lack of a rapid velocity decline at
early times appears to be a consistent trait of 03fg-like SNe, as shown in the
top panel of Fig. 12, indicating that the Si layer is confined to a narrow
range. ASASSN-15hy shows a velocity plateau at $\sim 8{,}600$ km s-1 from $-$5
to +2 d. Similar behavior was also reported in the bright 03fg-like object SN
2007if (Scalzo et al., 2010). In general, the velocities of ASASSN-15hy and
other 03fg-like SNe are relatively lower than the mean velocity of SS SNe Ia
from Folatelli et al. (2013), however, there is a large spread of velocity
among the group.

In ASASSN-15hy, it is not possible to determine the Si velocity evolution
after +2 d, because the blending of the low-excitation Fe ii $\lambda$6456,
$\lambda$6516 lines with the Si ii $\lambda$6335 feature causes a sudden drop
in in what is measured as the Si velocity. Therefore, after +2 d the
measurement does not represent the true velocity of the Si ii $\lambda$6355
feature. This blending effect is usually seen in subluminous SNe Ia at similar
early phases (Galbany et al., 2019) and at later times in normal SNe Ia (e.g.,
Folatelli et al., 2013). This velocity drop has been previously mistaken as a
distinguishing feature of 03fg-like objects (Scalzo et al., 2012). The fact
that we see this blending effect so early in ASASSN-15hy may be an indication
that there is a low ionization state in the line-forming region.

Figure 11:  Branch diagram (Branch et al., 2006) with the pseudo-equivalent
width of Si ii $\lambda$6355 and $\lambda$5972, The open markers are the
measurements from Folatelli et al. (2013), the filled diamonds with purple
edges are the measurements from this work of the currently available 03fg-like
objects.  Figure 12:  The time evolution of the Si ii $\lambda$6355 and C ii
$\lambda$6580 velocities (top panel) and the C ii $\lambda$6580 pseudo-
equivalent widths (bottom panel). ASASSN-15hy is shown compared to other 03fg-
like events and the normal SN 2011fe. The mean Si ii $\lambda$6355 velocity
measurements of the four Branch diagram subgroups from Folatelli et al. (2013)
are plotted in the top panel for comparison as well.

The C ii $\lambda$6580 velocity of ASASSN-15hy is similar to the Si ii
$\lambda$6355 velocity at $-$11 d ($\sim 9{,}300$ km s-1), but diverges from
the Si ii velocity shortly after this time period. By $B$-band maximum, the C
ii velocity is $\sim 2{,}500$ km s-1 lower than the Si iivelocity. This is in
contrast with the majority of the comparison SNe which show similar velocity
gradients for C ii $\lambda$6580 and Si ii $\lambda$6355.

Previous studies have shown that for normal SNe Ia the C ii and Si ii velocity
evolution is similar (Parrent et al., 2011; Folatelli et al., 2012). However,
ASASSN-15hy has a pre-maximum C ii velocity decline rate of $\sim 3{,}000$ km
s-1 per 10 days but only $\sim$500 km s-1 per 10 days for Si ii. This
diverging behavior in velocities is also seen in SN 2012dn (Chakradhari et
al., 2014; Parrent et al., 2016; Taubenberger et al., 2019). This may indicate
that there is carbon-rich material below the Si layer, or more likely, it
could be a projected velocity and ionization effect. If the C ii is in a
confined region just above the photosphere, more absorption is seen from the
material not moving directly towards the observer, which results in a
projection effect that makes the material appear to be at a lower velocity
(Hoeflich, 1990).

Despite the differences in velocity of C ii $\lambda$6580 within the 03fg-like
subgroup, ASASSN-15hy and most other members of this subgroup have large pEW
values for this ion, which continues past maximum light. It persists much
longer than in normal SNe Ia, such as SN 2011fe, as shown in the bottom panel
of Fig. 12. The C ii pEW of ASASSN-15hy has a value of $\sim$19 Å at $-$11 d
and gradually decreases with time, disappearing around 2 weeks past maximum
light. In contrast, the normal SN Ia 2011fe only has a pEW of $\sim$5 Å at
$-$17 d, which rapidly declines to $\sim$2 Å within 5 days, and then gradually
disappears while approaching maximum light. This is evidence of substantial
unburned material within the ejecta of ASASSN-15hy.

###  5.2 Near-infrared Wavelengths

(a) NIR spectra of ASASSN-15hy. (b) Velocity measurements of assumed ions.

Figure 13: The line identifications of the NIR spectra of ASASSN-15hy. Panel
(a) shows the first five NIR spectra of ASASSN-15hy and the most likely ion
labels for the corresponding features. The vertical gray bands mark the
regions of strong telluric adsorptions. Panel (b) presents the velocity
measurements of several assumed ions. The open markers represent features with
possible blending. The velocities of optical Si ii and C ii are also plotted
as reference. The colors of the features are consistent in the two panels.
Figure 14: NIR spectral comparison of ASASSN-15hy with three SNe Ia: the 03fg-
like object SN 2009dc (Taubenberger et al., 2011) (left panel), the normal SN
2012fr (Contreras et al., 2018; Hsiao et al., 2019) (middle panel), and the
transitional iPTF13ebh (Hsiao et al., 2015) (right panel). The vertical red
lines in the right panel mark the positions of prominent C i lines in
iPTF13ebh. The regions of the strongest telluric adsorptions are marked with
vertical gray bands.

ASASSN-15hy has 6 NIR spectra covering epochs from $-$8.5 to +71.7 d, making
it only the second 03fg-like object with NIR spectral observations (the first
one being SN 2009dc). All spectra have high S/N allowing for unambiguous
identifications of several lines and their evolution.

####  5.2.1 Line Identification and Evolution

A number of spectral lines in ASASSN-15hy are identified and labeled in Fig.
13 (a). The line list from Marion et al. (2009), as well as line
identifications from Taubenberger et al. (2011), Hsiao et al. (2015, 2019) and
Gall et al. (2012), were used as a guide to identify the NIR features in
ASASSN-15hy. The NIR spectra of ASASSN-15hy mainly consist of Ca ii, Mg ii, Fe
ii, and Co ii lines with a possible contribution from C i.

The strongest feature in the NIR is the Ca ii 0.8538 $\mu$m IR triplet, which
is also present in the region overlapping with the optical spectra. Multiple
Mg ii lines also were identified: 0.9227, 1.0927, 1.6787, and 2.1369 $\mu$m.
These line identifications were confirmed by examining these features in
velocity space, where they all follow a similar velocity trend, shown as solid
symbols in the left panel of Fig. 13 (b). The Mg ii 0.9227 $\mu$m feature may
also have contributions from O i and Si ii, but in ASASSN-15hy these are minor
as the velocity evolution of this feature matches the other Mg ii lines. We
attribute the absorption feature seen at $\sim$0.97 $\mu$m (marked orange in
Fig. 13) to be a blend of Mg ii 1.0092 $\mu$m and Fe ii 0.9998 $\mu$m. Thus,
the blending likely affects the velocity evolution of these two ions presented
in Fig. 13 (b). This feature is likely to have a significant contribution from
Fe ii by +8.1 d. The presence of Fe ii at this time period is consistent with
the optical spectroscopy of ASASSN-15hy, where a velocity drop in the Si ii
$\lambda$6355 feature due to the blending with Fe ii $\lambda$6456 and
$\lambda$6516 was observed around the same time.

The IGE features become stronger with time. As shown in the +22.8 d spectra,
Fe ii 0.9998, 1.0500, and 1.0863 $\mu$m are all clearly visible and are also
seen in SN 2009dc (Taubenberger et al., 2011). Co ii 2.1365, 2.2205, 2.3613,
and 2.4596 $\mu$m emission features are also present in the 2.1 $-$ 2.4 $\mu$m
region in the same spectrum.

####  5.2.2 Possible C i Detection

Throughout the NIR spectral observations, there are no strong C i features
such as those observed in SN 1999by (Hoeflich et al., 2002), iPTF13ebh (Hsiao
et al., 2015) or SN 2015bp (Wyatt et al., 2021), but a weak C i 1.0693 $\mu$m
line may be present in ASASSN-15hy at $-$8.5 and $-$1.5 d. the right panel of
Fig. 14 shows a notch at $\sim$1.03 $\mu$m (marked with dark red) next to the
absorption of Mg ii 1.0927 $\mu$m. This is similar to the position of the C i
1.0693 $\mu$m feature in iPTF13ebh at early times, where the feature is strong
and unambiguous. The weak notch in ASASSN-15hy would correspond to a velocity
of $\sim 12{,}000$ km s-1 assuming that it originated from C i 1.0693 $\mu$m.
Such a velocity at $-1.5$ d is $\sim 6{,}000$ km s-1 faster than the optical C
ii $\lambda$6580 line. This velocity difference could be due to a large carbon
envelope in the ejecta that extends to higher velocities. An ionization change
throughout the envelope could cause C i and C ii to have distinctly different
observed velocities. Furthermore, if the C ii layer is small, the velocity
could appear lower than the location of the bulk of the C ii material due to
projected velocity effects (see Section 5.1.2).

Unfortunately, there are no other clear C i features, such as C i 1.1754
$\mu$m, in the first two NIR spectra. The next observation at +8.1 d contains
two newly formed features aroun 1.02 and 1.12 $\mu$m, marked red and pink in
Fig. 13 (a), respectively. These may be attributed to C i 1.0693 $\mu$m and
1.1754 $\mu$m. However, these lines are more likely produced through
absorption from IGEs, due to their similarity in velocity to the Fe ii lines.
We do not exclude the possibility that these two features are blended with
weak C i.

####  5.2.3 Missing H-band Break

Interestingly, ASASSN-15hy does not show an $H$-band break near 1.5 $\mu$m
until +72 d. For a normal SN Ia, the $H$-band break (Kirshner et al., 1973)
appears around +3 d and peaks in strength around $\sim$+10 to +12 d (e.g.,
Hsiao et al., 2013), as demonstrated by SN 2012fr (Contreras et al., 2018;
Hsiao et al., 2019) in the middle panel of Fig. 14. The $H$-band break is
produced by a multiplet of allowed Co ii, Ni ii, and Fe ii emission lines that
emerge when the photosphere recedes deep into the 56Ni rich region (Wheeler et
al., 1998; Hoeflich et al., 2002). The strength of the $H$-band break has been
found to correlate with SN Ia light-curve shape (Hsiao et al., 2013), and the
outer blue velocity ($v_{\rm{edge}}$) of the break has been found to directly
trace the outer 56Ni region in the ejecta (Ashall et al., 2019a, b). The
absence of an $H$-band break suggests that the photosphere has not reached the
core of the 56Ni region by +30 d. This absence was also noted in SN 2009dc
(Taubenberger et al., 2011).

Despite the peculiar behavior of the $H$-band break, Co ii features are
apparent in the 2.1$-$2.4 $\mu$m region at +23 d, which is a similar phase to
the onset of these features in all three comparison SNe in Fig. 14. It has
been proposed that these Co ii features form in the same region as the
$H$-band break (Wheeler et al., 1998). We rule out the lack of the $H$-band
break being caused by an ionization effect, since 91T-like SNe Ia show strong
$H$-band breaks but have a higher-ionization state than ASASSN-15hy (M.M.
Phillips et al. in preparation). The missing $H$-band break in ASASSN-15hy
could be evidence of an opacity difference between the $H$ and $K$ bands for
this peculiar object (see Appendix D.3 for more details).

The NIR spectra of ASASSN-15hy and SN 2009dc are nearly identical in both
color and evolution of the line profiles, as shown in the left panel of Fig.
14. ASASSN-15hy shows both similarities and discrepancies when compared to a
normal SN Ia. The early time NIR spectra of ASASSN-15hy have smooth continua
and show shallow IME features such as Mg ii, which is consistent with both SN
2009dc and SN 2012fr. However, after maximum light the features of ASASSN-15hy
are different from the normal SN Ia 2012fr, especially in the 1.2 to 1.6
$\mu$m region. As discussed above, part of this is due to the delayed onset of
the $H$-band break and distribution of 56Ni. After maximum light, SN 2012fr
presents strong absorption features near 1.2 $\mu$m, thought to be the cause
of the rapid decrease in $J$-band magnitude that is seen in normal SNe Ia
(Branch et al., 1983; Elias et al., 1985). However, ASASSN-15hy shows a
limited decrease of the flux in this region, which is consistent with its
broad and slowly-declining $J$-band light curve (see Fig. 7).

##  6 Model Results

In this section, we briefly discuss the physics of ASASSN-15hy, provide one
possible interpretation of the data, and compare the data to explosion models.
A more in depth description of our modeling assumptions and methodology can be
found in Appendix D.

###  6.1 Envelope models

In this work, we employ the parameterized framework of spherical envelope
models (Hoeflich & Khokhlov, 1996) and make use of the observed spectral
evolution and photometric properties of ASASSN-15hy to constrain model
parameters. This class of models has been previously shown to provide a
consistent picture for the 03fg-like object LSQ14fmg (Hsiao et al., 2020) and
the observed slow rise and decline of SN 2009dc (Noebauer et al., 2016). The
envelope models may be consistent with the core degenerate (CD) scenario
(e.g., Kashi & Soker, 2011). The CD scenario consists of the explosion of a
degenerate core within a nondegenerate envelope. In this scenario, the
ignition process and thermonuclear explosion may be triggered by the merger
between the core of an AGB star and a WD: 1) on dynamical time scale which
results in a detonation (Kashi & Soker, 2011; Aznar-Siguán et al., 2015;
Taubenberger, 2017), or 2) on secular time scales which may lead to an early
deflagration phase with a transition to a detonation, subsequently called
“Deflagration-CD” (DCD) (Hoeflich et al., 2019). We stress that in our models
the nature of the envelope is unknown. It may be the nondegenerate outer
envelope of an AGB star in the CD scenario, but we do not rule out other
possibilities.

The main model parameters are: the mass of the nondegenerate envelope
$M_{\text{env}}$; the radius of the nondegenerate envelope $R_{\text{env}}$;
the extent of the He and C layers in mass and velocity space; the initial
metallicity $Z$; the mass of the hydrostatic, possibly rotating, core which is
referred to as the core mass, $M_{\text{core}}$; the amount of mass burnt
during the deflagration phase, $M_{\text{defl}}$, which is controlled by the
transition density at which the deflagration transforms into a detonation
$\rho_{\text{tr}}$; and possible interaction with the nearby environment or
wind.

The central density ($\rho_{c}$) is calculated using the equations of
hydrostatic equilibrium for models with $M_{\text{core}}$ up to
$M_{\text{Ch}}$ (Hoeflich et al., 1998). For models with $M_{\text{core}}$ in
excess of $M_{\text{Ch}}$, we assume fast rotating cores that ignite at a
central density of $7\times 10^{9}$ g cm-3 (Yoon & Langer, 2005). These
densities are found at the lower end of the classical delayed-detonation
models and may be attributed to He- or C-accreting WDs within the
$M_{\text{Ch}}$ class of explosions (Telesco et al., 2015; Diamond et al.,
2015, 2018). For these values of $\rho_{c}$, approximately $0.08$ $M_{\odot}$
of electron rich iron-group elements are produced.

To first order, specific observational parameters are directly linked to model
parameters. This produces a set of selection criteria that allows us to
determine the best matching model. The main criteria are: $M_{\text{env}}$,
which determines the final “shell” velocity indicated by the Si line region
formed in quasi statistical equilibrium (QSE) (Hoeflich & Khokhlov, 1996;
Quimby et al., 2007); and $R_{\text{env}}$, which determines the width of the
shell. With increasing radius, the shell becomes more confined in velocity
space (Fig. D19); and $M_{\text{core}}$ determines the diffusion time scales
and therefore the rise time to maximum light (Hoeflich & Khokhlov, 1996;
Dessart et al., 2014; Shen & Moore, 2014). The detonation in a
sub-$M_{\text{Ch}}$ ($M_{\text{core}}\sim$1 $M_{\odot}$) model produces a rise
time that is too short, and models with $M_{\text{core}}=1.8~{}M_{\odot}$
produce a rise time that is too long by several days. The best agreement was
found with a model of $M_{\text{core}}=1.47~{}M_{\odot}$ that puts ASASSN-15hy
near, but over $M_{\text{Ch}}$.

Figure 15: Comparison of the absolute $BVi$-band light curves between models
and ASASSN-15hy. The light curves are shown relative to explosion and $B$-band
maximum. The observations have been $K$-corrected (except for the nondetection
limits in open squares) and corrected for Galactic reddening. Left panel: The
low-mass high-metallicity DET2ENV6 model ($M_{\rm{core}}=1.2~{}M_{\odot}$ and
$Z=Z_{\odot}$) from Hoeflich & Khokhlov (1996) is presented as an illustrative
case. Right panel: The best matching DCD07 model with $Z=10^{-4}~{}Z_{\odot}$
(solid lines) is shown in comparison to the DCD07 model with
$Z=10^{-10}~{}Z_{\odot}$ (dashed lines).

The secondary selection criteria in the case of DCD are: $\rho_{\text{tr}}$ or
$M_{\text{defl}}$, which regulates the 56Ni production and therefore the
luminosity, if all other parameters are kept the same, in the same fashion as
in the classical delayed detonation models (e.g., Hoeflich et al., 2002). In
general within the CD scenario, the metallicity, $Z$ and the relative amount
of $M_{\text{He}}$ control the color $B-V$, the UV and NIR flux, and the
formation of CO.

###  6.2 Best Matching Model

ASASSN-15hy’s relatively low peak luminosity and broad light curves place it
into a specific, and somewhat unusual, area of parameter space. As ASASSN-15hy
has a low Si velocity at maximum light, which translates into a relative large
$M_{\text{env}}$, previous sets of envelope models which detonate (Khokhlov et
al., 1993; Hoeflich & Khokhlov, 1996; Quimby et al., 2007) predict a higher
peak luminosity than majority of other 03fg-like events and generally slower
rise and decline (Noebauer et al., 2016). The secondary model criteria such as
$M_{\text{defl}}$ and $Z$ may be the keys to explain the peculiarity of
ASASSN-15hy. The explanation under the DCD envelope model, the underlying
physics, and the path to arrive at the best matching parameters are presented
in Appendix D. Here we summarize the main results.

The best matching model (a slightly modified version of DCD07, see Appendix
D.4) has the following parameters: $M_{\text{env}}=0.7~{}M_{\odot}$,
$R_{\text{env}}=10~{}R_{\text{core}}$, $M_{\text{He,env}}=0.1~{}M_{\odot}$,
$M_{\text{core}}=1.47~{}M_{\odot}$, $Z=10^{-4}~{}Z_{\odot}$,
$M_{\text{defl}}=0.42~{}M_{\odot}$, and a total 56Ni mass $=$
$0.87~{}M_{\odot}$. The large $M_{\text{core}}$ was required to produce long
diffusion time scales and a broad light curve; a large $M_{\text{env}}$ is
required to produce the low observed ejecta velocities exemplified by Si ii;
the low metallicity produces the long rise and lack of line blanketing
relative to the normal SNe Ia; and the low transition density or large
$M_{\text{defl}}$ ensures that the 56Ni mass is lower than that which is
produced in other 03fg-like objects and thus, the luminosity of the SN is low.
Finally, a small $R_{\text{env}}$, which was determined by the velocity range
of Si ii, produces an extended density enhancement in the outer layers of the
ejecta rather than a confined density shell, which would be produced by larger
envelope radii. The chemical structure is largely consistent with the observed
velocity and velocity range of the ions (see Appendix D.2 and D.3). Note that
up to 40% of the total luminosity comes in the form of hard X-rays and UV at
early times in our model, as discussed in Appendix D.4.

In Fig. 15, the effects of metallicity $Z$ and core mass $M_{\text{core}}$ are
illustrated using the following models: the high $Z$ ($Z=Z_{\odot}$) and low
$M_{\text{core}}$ ($M_{\text{core}}=1.2~{}M_{\odot}$) DET2ENV6 model of
Hoeflich & Khokhlov (1996) in the left panel and our best matching DCD07 with
$Z=10^{-4}~{}Z_{\odot}$ in comparison to DCD07 with $Z=10^{-10}~{}Z_{\odot}$
in the right panel. The low-mass DET2ENV6 model has $0.6~{}M_{\odot}$ of
$M_{\text{env}}$, similar to the best matching parameter of ASASSN-15hy, but
shows a rise time significantly shorter than the observation and a very
asymmetric light curve. Moreover, its high $Z$ leads to an optically thick C/O
shell, and the low $M_{\text{core}}$ leads to a shorter rise time and faster
evolution that is inconsistent with ASASSN-15hy. The DCD07 models
($M_{\text{core}}=1.47~{}M_{\odot}$), as shown in the right panel of Fig. 15,
provide overall matching light-curve shapes and luminosities as the observed
ones. Comparing the $Z=10^{-4}~{}Z_{\odot}$ and $Z=10^{-10}~{}Z_{\odot}$ DCD07
models shows a more extended dark phase as $Z$ increases, and the
$Z=10^{-4}~{}Z_{\odot}$ model is more consistent with the early ASAS-SN
nondetections.

ASASSN-15hy requires a rather large envelope mass that has a significant
amount of He left in the outer layers combined with a low metallicity and a
late deflagration to detonation transition. The significant amount of helium
and the low progenitor metallicity below that of the host galaxy points to an
old, low-mass progenitor. Note that, despite the considerable amount of He,
narrow He lines may not be detected due to the low opacity in the outer layer
and gamma-ray trapping, such that the necessary gamma-rays to excite the He
are not available (e.g., Graham, 1988). ASASSN-15hy is not a typical 03fg-like
object compared to the current public sample. An analysis on larger sample is
needed to determine whether the model parameter space of ASASSN-15hy is common
for 03fg-like SNe or not.

###  6.3 DCD Models in Context

In the following discussion, we put ASASSN-15hy and our model in a larger
context. The basic characteristics are the explosion of a degenerate core of
$\approx M_{\text{Ch}}$ embedded in a compact high-mass envelope with a
significant fraction of He $\approx 0.1-0.2~{}M_{\odot}$ and low $Z$. The
resulting explosion has a red color at maximum light for a 03fg-like object, a
low luminosity, and took place in a low-metallicity galaxy with ongoing star
formation.

The best matching model has a low transition density from deflagration to
detonation (DDT) which results in an extend phase of the deflagration burning
similar to the sublumious classical SNe Ia models. As discussed in Appendix
D.4 a low metallicity and C/O ratio can be expected to shift the DDT to low
densities (Poludnenko et al., 2019). Thus, the very low Z may be the key for
understanding why ASASSN-15hy is a subluminous CD explosion.

Our analysis suggests that ASASSN-15hy can be described by the DCD scenario in
which the explosion starts as a rather long deflagration phase and turns into
a detonation. An initial deflagration requires a WD companion merging with the
central core of an AGB star on secular time scales rather than on dynamical
ones. This is similar to classical delayed detonation $M_{\text{Ch}}$ models,
which have a long deflagration phase and low transition densities. These
models may explain transitional and subluminous SNe Ia (Hoeflich et al., 2002;
Patat et al., 2012; Hsiao et al., 2015; Ashall et al., 2016, 2018; Galbany et
al., 2019).

In normal SNe Ia, a low transition density is rare and produces intrinsically
red optical colors similar to that observed in ASASSN-15hy. Within the DCD
framework, subluminous objects that behave like ASASSN-15hy should be rare and
the low metallicity suggests that the progenitor may have been a Pop II/III
star. Low-mass AGB stars in metal-poor environments are carbon stars with
highly enhanced Ba and other s-process elements (Aaronson & Mould, 1980; Cohen
et al., 2006; Kirby et al., 2015). Even though hydrogen is not apparent in the
spectra and the limits for ongoing mass loss are rather low, hydrogen and
s-process elements may be detectable when the ejecta have cleared the low-
density cocoon that has a size of a few tenths to a few light years and was
produced by prior mass-loss episodes (Dragulin & Hoeflich, 2016). We do not
expect Balmer lines to appear right away, but when the envelope collides with
the edge of the cocoon in $\sim 1-10$ yr. This collision may lead to a revival
of X-ray emission.

For degenerate cores (or WDs), the C/O ratio increases with decreasing main
sequence mass $M_{\text{MS}}$ as a result of stellar evolution. At lower mass
and hence lower central temperatures, during the central stellar He burning,
$\alpha$ capture on 12C dominates triple-$\alpha$, resulting in low C/O ratios
(C/O $\sim 0.1-0.2$). The subsequent He-shell burning occurs at higher
temperatures. Thus, shell burning results in ratios close to the statistical
equilibrium (C/O $\approx 1$). Because of the size of the convective He
burning core decreases with $M_{\text{MS}}$ and $Z$ (Domínguez et al., 2001),
the overall C/O ratio also decreases with $M_{\text{MS}}$ and $Z$. The CD
scenario requires an explosion during the AGB phase. A common-envelope phase
is more common for more massive stars. Stars of mass $\sim 5-7~{}M_{\odot}$
have evolutionary times of $\sim 10^{8}$ years and may be expected to be
related to starburst episodes. Although there is a strong correlation between
low-metallicity, starburst galaxies and 03fg-like in general (L. Galbany et
al, in preparation), the progenitor system of ASASSN-15hy may have been
produced in a previous episode of star formation. This may indicate that 03fg-
like SNe come from a diverse group of progenitors.

###  6.4 Alternative Scenarios

Under the class of interaction or envelope models, several studies on 03fg-
like objects have suggested that the origin of the envelope material could
come from WD-WD mergers (e.g., Scalzo et al., 2010; Hachinger et al., 2012;
Taubenberger et al., 2013; Noebauer et al., 2016). Here we briefly discuss the
alternative progenitor systems based on the parameter space described above.
However, numerical simulations and comparisons are beyond the scope of this
work.

The dynamical merger of two WDs may be an alternative possible scenario for
ASASSN-15hy and 03fg-like objects (e.g., Pakmor et al., 2010, 2012). This
scenario may produce a similar overall structure, light curves, and spectra
when compared to CDs of $\sim 2~{}M_{\odot}$ that undergo detonation burning.
However, it is not obvious how, in this framework, one would produce an
explosion with similar conditions as ASASSN-15hy. Our models of ASASSN-15hy
require an extended nondegenerate envelope consisting of He/C/O without
extensive mixing and high density outer material, which favors a CD/AGB
scenario. Furthermore, a long deflagration phase is required to produce the
low 56Ni mass, which cannot be achieved in a dynamical merger, since these
models detonate. Finally a total ejecta mass exceeding the $M_{\text{Ch}}$ is
required to produce the observed light curves. This may disfavor the secular
merger scenario (Piersanti et al., 2003), which explodes near $M_{\text{Ch}}$,
as a viable scenario for ASASSN-15hy.

Though there is no polarimetric observations available for ASASSN-15hy, the
lack of polarization in the other two 03fg-like explosions, SN 2009dc (Tanaka
et al., 2010) and SN 2007if (Cikota et al., 2019), may disfavor explosion
mechanisms with strong, large scale asymmetries, such as dynamical mergers
(Bulla et al., 2016). However, this does not exclude the possibility that a
majority of 03fg-like SNe are the result of dynamical mergers. ASASSN-15hy and
LSQ14fmg (Hsiao et al., 2020) are consistent with the DCD scenario, however,
these two are on the extreme ends among the current 03fg-like objects. A
larger sample of 03fg-like events and larger scale of model comparison are
needed for future analysis.

##  7 Conclusion

Photometric and spectroscopic observations of ASASSN-15hy have been presented,
covering $-$12 to +155 d relative to $B$-band maximum. ASASSN-15hy has many
similar characteristics to previously discovered 03fg-like SNe Ia. However,
ASASSN-15hy also shows its own uniqueness and provides additional insight into
this peculiar subgroup.

ASASSN-15hy has the following properties that are similar to other 03fg-like
SNe:

  * •

Slowly evolving (rise time = 22.5 $\pm$ 4.6 d in $V$-band) and broad light
curves ($\Delta{\rm{m}_{15}}(B)$ = 0.72 $\pm$ 0.04 mag),

  * •

No distinct secondary maxima and peaking later than normal SNe Ia in the
$iYJH$ bands,

  * •

A relatively blue UV$-$optical color ($uvm2-u$ = $1.72\pm 0.16$ mag at
$B$-band maximum),

  * •

Bright peak luminosity in both UV (M${}_{uvm2}=$ $-18.17^{+0.11}_{-0.27}$ mag)
and NIR ($M_{H}\leq$ $-$19.39 mag),

  * •

Weak Ca ii but strong and persistent C ii spectral feature,

  * •

A relatively low Si ii velocity ($V_{\text{Si}}=8{,}600\pm 200$ km s-1 at
maximum light) and lacks an early time rapid velocity decline,

  * •

No prominent $H$-band break feature in the first month past $B$-band maximum.

AMUSING observations show that the host of ASASSN-15hy is a low-mass, low-
metallicity galaxy with a relatively young stellar population. These
properties are generally consistent with other 03fg-like events. The host
galaxy of ASASSN-15hy is barely visible in optical images, but in the
H$\alpha$ map, the host is clearly visible.

Despite similarities mentioned above with previously published 03fg-like SNe,
ASASSN-15hy is unique in several ways. At optical wavelengths, it is dim
compared with classical 03fg-like events and even when compared with normal
SNe Ia with similar decline rates. It has a peak absolute magnitude of
M${}_{B}=$ $-19.14^{+0.11}_{-0.16}$ mag, which is a $\sim 1$ mag dimmer than
SN 2009dc and located below the LWR of normal SNe Ia. This makes ASASSN-15hy
the faintest SN among the current 03fg-like SNe Ia sample. ASASSN-15hy is also
the reddest and has a observed $B-V$ color of 0.18 $\pm$ 0.01 mag at $B$-band
maximum, which is $\sim$0.1 mag redder than the average of other 03fg-like
objects. Furthermore, ASASSN-15hy has the longest delayed onset of its
$i$-band primary peak relative to that in the $B$ band at 7.2 $\pm$ 1.1 d.

We find that ASASSN-15hy is consistent with the “deflagration core-degenerate”
scenario, which consists of the secular merger of the degenerate core of an
AGB star and a WD inside the nondegenerate envelope of the AGB star. The best
matching model has the parameters: $M_{\text{env}}=0.7~{}M_{\odot}$,
$R_{\text{env}}=10~{}R_{\text{core}}$, $M_{\text{He,env}}=0.1~{}M_{\odot}$,
$M_{\text{core}}=1.47~{}M_{\odot}$, $Z=10^{-4}~{}Z_{\odot}$,
$M_{\text{defl}}=0.42~{}M_{\odot}$, and total
$M(^{56}\rm{Ni})=0.87~{}M_{\odot}$. A large core mass is required in order to
produce a broad light curve. To produce a slow expansion velocity and a narrow
IME shell, a large envelope mass is needed. A major portion of this envelope
consists of He, which does not add to the line opacity but contributes to the
hydrodynamics in the ejecta. Critically, a long deflagration phase is required
to produce the observed low luminosity and red $B-V$ color. Finally, a
metallicity of $Z\sim 10^{-4}~{}Z_{\odot}$ is required to provide a small
opacity and a dark phase on the rise, in order to match the early light curve
and avoid producing significant line-blanketing to block the UV flux and over-
redistribute the flux to the NIR.

The predominant host-galaxy type for 03fg-like objects has been low mass, low
metallicity and star-forming, including ASASSN-15hy, meaning that these
objects may be more common at high redshifts. This has direct consequences for
high-redshift dark energy experiments. 03fg-like SNe Ia have similar light-
curve shapes in the $B$ and $V$ bands compared to normal SNe Ia and therefore
cannot be distinguished at high redshifts in these rest-frame bands alone.
Observations of redder bands in the rest frame, such as $i$ band to NIR band,
are required to identify 03fg-like SNe Ia through purely photometric means.
Further studies on larger samples of 03fg-like SNe Ia are necessary to fully
understand their impact on dark energy experiments. Overall, more observations
are required to test the consistency of the model simulations to understand
the physics of the progenitors and the explosions.

## Acknowledgments

The authors would like to thank the anonymous referee for their useful
comments. We thank the technical, scientific staff, and support astronomers of
the Las Campanas Observatory and the ESO Telescopes at the Paranal
Observatory. The CSP-II has been supported by NSF grants AST-1008343,
AST-1613426, AST-1613455, AST-1613472, and the Danish Agency for Science and
Technology and Innovation through a Sapere Aude Level 2 grant (PI: M.S.) C.A.
is supported by the NSF grant AST #1908952. J.L., S.K., M.S., and E.Y.H
acknowledge the support provided by the Florida Space Research Program. P.A.H.
acknowledges the support by the National Science Foundation grant AST-1715133.
L.G. acknowledges financial support from the Spanish Ministry of Science,
Innovation and Universities (MICIU) under the 2019 Ramón y Cajal program
RYC2019-027683 and from the Spanish MICIU project PID2020-115253GA-I00. E.B.
was supported in part by NASA grant 80NSSC20K0538. M.D.S. is supported by
generous grants from Villum FONDEN (13261, 28021) and by a project grant
(8021-00170B) from the Independent Research Fund Denmark. Support for
T.W.-S.H. was provided by NASA through the NASA Hubble Fellowship grant HST-
HF2-51458.001-A awarded by the Space Telescope Science Institute, which is
operated by the Association of Universities for Research in Astronomy, Inc.,
for NASA, under contract NAS5-26555. H.K. was funded by the Academy of Finland
projects 324504 and 328898. Support for J.L.P. is provided in part by ANID
through the Fondecyt regular grant 1191038 and through the Millennium Science
Initiative grant ICN12$\\_$009, awarded to The Millennium Institute of
Astrophysics, MAS. Time domain research by D.J.S. is supported by NSF grants
AST-1821987, 1813466, & 1908972, and by the Heising-Simons Foundation under
grant #2020-1864. This paper includes data gathered with the 6.5 meter
Magellan Telescopes located at Las Campanas Observatory, Chile, and is also
based upon observations made with ESO Telescopes at the La Silla or Paranal
Observatories under programme ID(s) 191.D-0935 and 099.D-0022(A). This work
has been partially supported by the Spanish grant PGC2018-095317-B-C21 within
the European Funds for Regional Development (FEDER). Finally, this paper is
also includes observations obtained at the international Gemini Observatory
(GN-2015A-Q-8, GS-2015A-Q-5), a program of NSF’s NOIRLab, which is managed by
the Association of Universities for Research in Astronomy (AURA) under a
cooperative agreement with the National Science Foundation. on behalf of the
Gemini Observatory partnership: the National Science Foundation (United
States), National Research Council (Canada), Agencia Nacional de Investigación
y Desarrollo (Chile), Ministerio de Ciencia, Tecnología e Innovación
(Argentina), Ministério da Ciência, Tecnologia, Inovações e Comunicações
(Brazil), and Korea Astronomy and Space Science Institute (Republic of Korea).

$Facilities$: Swope (e2v), du Pont (RetroCam, WFCCD), $Swift$, ASAS-SN, VLT
(MUSE), NOT (ALFOSC), LT (SPART), FLWO 1.5m (FAST), Baade (FIRE), GN (GNIRS),
GS (F2).

$Software$: $SNooPy$ (Burns et al., 2011); HYDRA (Hoeflich, 2003, 2009).

## References

  * Aaronson & Mould (1980) Aaronson, M., & Mould, J. 1980, ApJ, 240, 804, doi: 10.1086/158293
  * Alard (2000) Alard, C. 2000, A&AS, 144, 363, doi: 10.1051/aas:2000214
  * Alard & Lupton (1998) Alard, C., & Lupton, R. H. 1998, ApJ, 503, 325, doi: 10.1086/305984
  * Arnett (1982) Arnett, W. D. 1982, ApJ, 253, 785, doi: 10.1086/159681
  * Ashall et al. (2016) Ashall, C., Mazzali, P. A., Pian, E., & James, P. A. 2016, MNRAS, 463, 1891, doi: 10.1093/mnras/stw2114
  * Ashall et al. (2018) Ashall, C., Mazzali, P. A., Stritzinger, M. D., et al. 2018, MNRAS, 477, 153, doi: 10.1093/mnras/sty632
  * Ashall et al. (2019a) Ashall, C., Hoeflich, P., Hsiao, E. Y., et al. 2019a, ApJ, 878, 86, doi: 10.3847/1538-4357/ab204b
  * Ashall et al. (2019b) Ashall, C., Hsiao, E. Y., Hoeflich, P., et al. 2019b, ApJ, 875, L14, doi: 10.3847/2041-8213/ab1654
  * Ashall et al. (2020) Ashall, C., Lu, J., Burns, C., et al. 2020, ApJ, 895, L3, doi: 10.3847/2041-8213/ab8e37
  * Ashall et al. (2021) Ashall, C., Lu, J., Hsiao, E. Y., et al. 2021, arXiv e-prints, arXiv:2106.12140.  https://arxiv.org/abs/2106.12140
  * Aznar-Siguán et al. (2015) Aznar-Siguán, G., García-Berro, E., Lorén-Aguilar, P., Soker, N., & Kashi, A. 2015, MNRAS, 450, 2948, doi: 10.1093/mnras/stv824
  * Bacon et al. (2010) Bacon, R., Accardo, M., Adjali, L., et al. 2010, in Society of Photo-Optical Instrumentation Engineers (SPIE) Conference Series, Vol. 7735, Ground-based and Airborne Instrumentation for Astronomy III, ed. I. S. McLean, S. K. Ramsay, & H. Takami, 773508, doi: 10.1117/12.856027
  * Baldwin et al. (1981) Baldwin, J. A., Phillips, M. M., & Terlevich, R. 1981, PASP, 93, 5, doi: 10.1086/130766
  * Blinnikov & Sorokina (2010) Blinnikov, S. I., & Sorokina, E. I. 2010, arXiv e-prints, arXiv:1009.4353.  https://arxiv.org/abs/1009.4353
  * Branch et al. (1983) Branch, D., Lacy, C. H., McCall, M. L., et al. 1983, ApJ, 270, 123, doi: 10.1086/161103
  * Branch et al. (2006) Branch, D., Dang, L. C., Hall, N., et al. 2006, PASP, 118, 560, doi: 10.1086/502778
  * Brown et al. (2014a) Brown, P. J., Breeveld, A. A., Holland, S., Kuin, P., & Pritchard, T. 2014a, Ap&SS, 354, 89, doi: 10.1007/s10509-014-2059-8
  * Brown et al. (2009) Brown, P. J., Holland, S. T., Immler, S., et al. 2009, AJ, 137, 4517, doi: 10.1088/0004-6256/137/5/4517
  * Brown et al. (2014b) Brown, P. J., Kuin, P., Scalzo, R., et al. 2014b, ApJ, 787, 29, doi: 10.1088/0004-637X/787/1/29
  * Bulla et al. (2016) Bulla, M., Sim, S. A., Pakmor, R., et al. 2016, MNRAS, 455, 1060, doi: 10.1093/mnras/stv2402
  * Burns et al. (2011) Burns, C. R., Stritzinger, M., Phillips, M. M., et al. 2011, AJ, 141, 19, doi: 10.1088/0004-6256/141/1/19
  * Burns et al. (2014) —. 2014, ApJ, 789, 32, doi: 10.1088/0004-637X/789/1/32
  * Cardelli et al. (1989) Cardelli, J. A., Clayton, G. C., & Mathis, J. S. 1989, ApJ, 345, 245, doi: 10.1086/167900
  * Chakradhari et al. (2014) Chakradhari, N. K., Sahu, D. K., Srivastav, S., & Anupama, G. C. 2014, MNRAS, 443, 1663, doi: 10.1093/mnras/stu1258
  * Chen et al. (2019) Chen, P., Dong, S., Katz, B., et al. 2019, ApJ, 880, 35, doi: 10.3847/1538-4357/ab2630
  * Childress et al. (2011) Childress, M., Aldering, G., Aragon, C., et al. 2011, ApJ, 733, 3, doi: 10.1088/0004-637X/733/1/3
  * Childress et al. (2016) Childress, M. J., Tucker, B. E., Yuan, F., et al. 2016, PASA, 33, e055, doi: 10.1017/pasa.2016.47
  * Cikota et al. (2019) Cikota, A., Patat, F., Wang, L., et al. 2019, MNRAS, 490, 578, doi: 10.1093/mnras/stz2322
  * Cohen et al. (2006) Cohen, J. G., McWilliam, A., Shectman, S., et al. 2006, AJ, 132, 137, doi: 10.1086/504597
  * Contreras et al. (2018) Contreras, C., Phillips, M. M., Burns, C. R., et al. 2018, ApJ, 859, 24, doi: 10.3847/1538-4357/aabaf8
  * Dessart et al. (2014) Dessart, L., Blondin, S., Hillier, D. J., & Khokhlov, A. 2014, MNRAS, 441, 532, doi: 10.1093/mnras/stu598
  * Diamond et al. (2015) Diamond, T. R., Hoeflich, P., & Gerardy, C. L. 2015, ApJ, 806, 107, doi: 10.1088/0004-637X/806/1/107
  * Diamond et al. (2018) Diamond, T. R., Hoeflich, P., Hsiao, E. Y., et al. 2018, ApJ, 861, 119, doi: 10.3847/1538-4357/aac434
  * Domínguez et al. (2001) Domínguez, I., Höflich, P., & Straniero, O. 2001, ApJ, 557, 279, doi: 10.1086/321661
  * Dopita et al. (2007) Dopita, M., Hart, J., McGregor, P., et al. 2007, Ap&SS, 310, 255, doi: 10.1007/s10509-007-9510-z
  * Dopita et al. (2010) Dopita, M., Rhee, J., Farage, C., et al. 2010, Ap&SS, 327, 245, doi: 10.1007/s10509-010-0335-9
  * Dopita et al. (2016) Dopita, M. A., Kewley, L. J., Sutherland, R. S., & Nicholls, D. C. 2016, Ap&SS, 361, 61, doi: 10.1007/s10509-016-2657-8
  * Dragulin & Hoeflich (2016) Dragulin, P., & Hoeflich, P. 2016, ApJ, 818, 26, doi: 10.3847/0004-637X/818/1/26
  * Eikenberry et al. (2008) Eikenberry, S., Elston, R., Raines, S. N., et al. 2008, in Proc. SPIE, Vol. 7014, Ground-based and Airborne Instrumentation for Astronomy II, 70140V, doi: 10.1117/12.788326
  * Elias et al. (1985) Elias, J. H., Matthews, K., Neugebauer, G., & Persson, S. E. 1985, ApJ, 296, 379, doi: 10.1086/163456
  * Elias et al. (1998) Elias, J. H., Vukobratovich, D., Andrew, J. R., et al. 1998, in Proc. SPIE, Vol. 3354, Infrared Astronomical Instrumentation, ed. A. M. Fowler, 555–565, doi: 10.1117/12.317281
  * Fabricant et al. (1998) Fabricant, D., Cheimets, P., Caldwell, N., & Geary, J. 1998, PASP, 110, 79, doi: 10.1086/316111
  * Filippenko et al. (1992a) Filippenko, A. V., Richmond, M. W., Matheson, T., et al. 1992a, ApJ, 384, L15, doi: 10.1086/186252
  * Filippenko et al. (1992b) Filippenko, A. V., Richmond, M. W., Branch, D., et al. 1992b, AJ, 104, 1543, doi: 10.1086/116339
  * Fink et al. (2018) Fink, M., Kromer, M., Hillebrandt, W., et al. 2018, A&A, 618, A124, doi: 10.1051/0004-6361/201833475
  * Fisher et al. (1999) Fisher, A., Branch, D., Hatano, K., & Baron, E. 1999, MNRAS, 304, 67, doi: 10.1046/j.1365-8711.1999.02299.x
  * Folatelli et al. (2012) Folatelli, G., Phillips, M. M., Morrell, N., et al. 2012, ApJ, 745, 74, doi: 10.1088/0004-637X/745/1/74
  * Folatelli et al. (2013) Folatelli, G., Morrell, N., Phillips, M. M., et al. 2013, ApJ, 773, 53, doi: 10.1088/0004-637X/773/1/53
  * Foley et al. (2010) Foley, R. J., Narayan, G., Challis, P. J., et al. 2010, ApJ, 708, 1748, doi: 10.1088/0004-637X/708/2/1748
  * Friedman et al. (2015) Friedman, A. S., Wood-Vasey, W. M., Marion, G. H., et al. 2015, ApJS, 220, 9, doi: 10.1088/0067-0049/220/1/9
  * Frohmaier et al. (2015) Frohmaier, C., Prajs, S., Dennefeld, M., et al. 2015, The Astronomer’s Telegram, 7452, 1
  * Fryer et al. (2010) Fryer, C. L., Ruiter, A. J., Belczynski, K., et al. 2010, ApJ, 725, 296, doi: 10.1088/0004-637X/725/1/296
  * Galbany et al. (2014) Galbany, L., Stanishev, V., Mourão, A. M., et al. 2014, A&A, 572, A38, doi: 10.1051/0004-6361/201424717
  * Galbany et al. (2016a) Galbany, L., Anderson, J. P., Rosales-Ortega, F. F., et al. 2016a, MNRAS, 455, 4087, doi: 10.1093/mnras/stv2620
  * Galbany et al. (2016b) Galbany, L., Stanishev, V., Mourão, A. M., et al. 2016b, A&A, 591, A48, doi: 10.1051/0004-6361/201528045
  * Galbany et al. (2019) Galbany, L., Ashall, C., Höflich, P., et al. 2019, A&A, 630, A76, doi: 10.1051/0004-6361/201935537
  * Gall et al. (2012) Gall, E. E. E., Taubenberger, S., Kromer, M., et al. 2012, MNRAS, 427, 994, doi: 10.1111/j.1365-2966.2012.21999.x
  * Gehrels et al. (2004) Gehrels, N., Chincarini, G., Giommi, P., et al. 2004, ApJ, 611, 1005, doi: 10.1086/422091
  * Gerardy et al. (2007) Gerardy, C. L., Meikle, W. P. S., Kotak, R., et al. 2007, ApJ, 661, 995, doi: 10.1086/516728
  * Graham (1988) Graham, J. R. 1988, ApJ, 335, L53, doi: 10.1086/185338
  * Hachinger et al. (2012) Hachinger, S., Mazzali, P. A., Taubenberger, S., et al. 2012, MNRAS, 427, 2057, doi: 10.1111/j.1365-2966.2012.22068.x
  * Hachisu et al. (2012) Hachisu, I., Kato, M., Saio, H., & Nomoto, K. 2012, ApJ, 744, 69, doi: 10.1088/0004-637X/744/1/69
  * Henden et al. (2015) Henden, A. A., Levine, S., Terrell, D., & Welch, D. L. 2015, in American Astronomical Society Meeting Abstracts, Vol. 225, American Astronomical Society Meeting Abstracts #225, 336.16
  * Hicken et al. (2007) Hicken, M., Garnavich, P. M., Prieto, J. L., et al. 2007, ApJ, 669, L17, doi: 10.1086/523301
  * Hicken et al. (2012) Hicken, M., Challis, P., Kirshner, R. P., et al. 2012, ApJS, 200, 12, doi: 10.1088/0067-0049/200/2/12
  * Hoeflich (1990) Hoeflich, P. 1990, A&A, 229, 191
  * Hoeflich (1995) —. 1995, ApJ, 443, 89, doi: 10.1086/175505
  * Hoeflich (2003) Hoeflich, P. 2003, in Astronomical Society of the Pacific Conference Series, Vol. 288, Stellar Atmosphere Modeling, ed. I. Hubeny, D. Mihalas, & K. Werner, 185
  * Hoeflich (2009) Hoeflich, P. 2009, in American Institute of Physics Conference Series, Vol. 1171, Recent Directions in Astrophysical Quantitative Spectroscopy and Radiation Hydrodynamics, ed. I. Hubeny, J. M. Stone, K. MacGregor, & K. Werner, 161–172, doi: 10.1063/1.3250057
  * Hoeflich et al. (2002) Hoeflich, P., Gerardy, C. L., Fesen, R. A., & Sakai, S. 2002, ApJ, 568, 791, doi: 10.1086/339063
  * Hoeflich & Khokhlov (1996) Hoeflich, P., & Khokhlov, A. 1996, ApJ, 457, 500, doi: 10.1086/176748
  * Hoeflich et al. (1995) Hoeflich, P., Khokhlov, A. M., & Wheeler, J. C. 1995, ApJ, 444, 831, doi: 10.1086/175656
  * Hoeflich et al. (1993) Hoeflich, P., Mueller, E., & Khokhlov, A. 1993, A&A, 268, 570
  * Hoeflich et al. (1998) Hoeflich, P., Wheeler, J. C., & Thielemann, F. K. 1998, ApJ, 495, 617, doi: 10.1086/305327
  * Hoeflich et al. (2017) Hoeflich, P., Hsiao, E. Y., Ashall, C., et al. 2017, ApJ, 846, 58, doi: 10.3847/1538-4357/aa84b2
  * Hoeflich et al. (2019) Hoeflich, P., Ashall, C., Fisher, A., et al. 2019, in Nuclei in the Cosmos XV, Vol. 219, 187–194, doi: 10.1007/978-3-030-13876-9_31
  * Höflich (2003) Höflich, P. 2003, in Astronomical Society of the Pacific Conference Series, Vol. 288, Stellar Atmosphere Modeling, ed. I. Hubeny, D. Mihalas, & K. Werner, 185
  * Holoien et al. (2015) Holoien, T. W. S., Stanek, K. Z., Kochanek, C. S., et al. 2015, The Astronomer’s Telegram, 7450, 1
  * Howell et al. (2006) Howell, D. A., Sullivan, M., Nugent, P. E., et al. 2006, Nature, 443, 308, doi: 10.1038/nature05103
  * Hoyle & Fowler (1960) Hoyle, F., & Fowler, W. A. 1960, ApJ, 132, 565, doi: 10.1086/146963
  * Hristov et al. (2018) Hristov, B., Collins, D. C., Hoeflich, P., Weatherford, C. A., & Diamond, T. R. 2018, ApJ, 858, 13, doi: 10.3847/1538-4357/aab7f2
  * Hsiao et al. (2007) Hsiao, E. Y., Conley, A., Howell, D. A., et al. 2007, ApJ, 663, 1187, doi: 10.1086/518232
  * Hsiao et al. (2013) Hsiao, E. Y., Marion, G. H., Phillips, M. M., et al. 2013, ApJ, 766, 72, doi: 10.1088/0004-637X/766/2/72
  * Hsiao et al. (2015) Hsiao, E. Y., Burns, C. R., Contreras, C., et al. 2015, A&A, 578, A9, doi: 10.1051/0004-6361/201425297
  * Hsiao et al. (2019) Hsiao, E. Y., Philips, M. M., Marion, G. H., et al. 2019, PASP, 131, 014002, doi: 10.1088/1538-3873/aae961
  * Hsiao et al. (2020) Hsiao, E. Y., Hoeflich, P., Ashall, C., et al. 2020, ApJ, 900, 140, doi: 10.3847/1538-4357/abaf4c
  * Iben & Tutukov (1984) Iben, Jr., I., & Tutukov, A. V. 1984, ApJS, 54, 335, doi: 10.1086/190932
  * Ilkov & Soker (2012) Ilkov, M., & Soker, N. 2012, MNRAS, 419, 1695, doi: 10.1111/j.1365-2966.2011.19833.x
  * Karp et al. (1977) Karp, A. H., Lasher, G., Chan, K. L., & Salpeter, E. E. 1977, ApJ, 214, 161, doi: 10.1086/155241
  * Kashi & Soker (2011) Kashi, A., & Soker, N. 2011, MNRAS, 417, 1466, doi: 10.1111/j.1365-2966.2011.19361.x
  * Katz & Dong (2012) Katz, B., & Dong, S. 2012, arXiv e-prints, arXiv:1211.4584.  https://arxiv.org/abs/1211.4584
  * Kewley et al. (2001) Kewley, L. J., Dopita, M. A., Sutherland, R. S., Heisler, C. A., & Trevena, J. 2001, ApJ, 556, 121, doi: 10.1086/321545
  * Khan et al. (2011) Khan, R., Stanek, K. Z., Stoll, R., & Prieto, J. L. 2011, ApJ, 737, L24, doi: 10.1088/2041-8205/737/1/L24
  * Khokhlov et al. (1993) Khokhlov, A., Mueller, E., & Hoeflich, P. 1993, Astronomy and Astrophysics, 270, 223
  * Khokhlov (1995) Khokhlov, A. M. 1995, ApJ, 449, 695, doi: 10.1086/176091
  * Khokhlov et al. (1997) Khokhlov, A. M., Oran, E. S., & Wheeler, J. C. 1997, ApJ, 478, 678, doi: 10.1086/303815
  * Kirby et al. (2015) Kirby, E. N., Guo, M., Zhang, A. J., et al. 2015, ApJ, 801, 125, doi: 10.1088/0004-637X/801/2/125
  * Kirshner et al. (1973) Kirshner, R. P., Willner, S. P., Becklin, E. E., Neugebauer, G., & Oke, J. B. 1973, ApJ, 180, L97, doi: 10.1086/181161
  * Kochanek et al. (2017) Kochanek, C. S., Shappee, B. J., Stanek, K. Z., et al. 2017, PASP, 129, 104502, doi: 10.1088/1538-3873/aa80d9
  * Krisciunas et al. (2017) Krisciunas, K., Contreras, C., Burns, C. R., et al. 2017, AJ, 154, 211, doi: 10.3847/1538-3881/aa8df0
  * Kuncarayakti et al. (2018) Kuncarayakti, H., Anderson, J. P., Galbany, L., et al. 2018, A&A, 613, A35, doi: 10.1051/0004-6361/201731923
  * Kushnir et al. (2013) Kushnir, D., Katz, B., Dong, S., Livne, E., & Fernández, R. 2013, ApJ, 778, L37, doi: 10.1088/2041-8205/778/2/L37
  * Lentz et al. (2000) Lentz, E. J., Baron, E., Branch, D., Hauschildt, P. H., & Nugent, P. E. 2000, ApJ, 530, 966, doi: 10.1086/308400
  * Lira (1996) Lira, P. 1996, Master’s thesis, Universidad de Chile
  * Livio & Riess (2003) Livio, M., & Riess, A. G. 2003, ApJ, 594, L93, doi: 10.1086/378765
  * Livne (1999) Livne, E. 1999, ApJ, 527, L97, doi: 10.1086/312405
  * Maeda et al. (2009) Maeda, K., Kawabata, K., Li, W., et al. 2009, ApJ, 690, 1745, doi: 10.1088/0004-637X/690/2/1745
  * Maoz et al. (2014) Maoz, D., Mannucci, F., & Nelemans, G. 2014, ARA&A, 52, 107, doi: 10.1146/annurev-astro-082812-141031
  * Marino et al. (2013) Marino, R. A., Rosales-Ortega, F. F., Sánchez, S. F., et al. 2013, A&A, 559, A114, doi: 10.1051/0004-6361/201321956
  * Marion et al. (2009) Marion, G. H., Höflich, P., Gerardy, C. L., et al. 2009, AJ, 138, 727, doi: 10.1088/0004-6256/138/3/727
  * Matheson et al. (2005) Matheson, T., Blondin, S., Foley, R. J., et al. 2005, AJ, 129, 2352, doi: 10.1086/429679
  * Mazzali et al. (2001) Mazzali, P. A., Nomoto, K., Cappellaro, E., et al. 2001, ApJ, 547, 988, doi: 10.1086/318428
  * Mazzali et al. (2007) Mazzali, P. A., Röpke, F. K., Benetti, S., & Hillebrandt, W. 2007, Science, 315, 825, doi: 10.1126/science.1136259
  * Mihalas (1978) Mihalas, D. 1978, Stellar Atmospheres, 2nd edn. (New York: W. H. Freeman)
  * Milne et al. (2013) Milne, P. A., Brown, P. J., Roming, P. W. A., Bufano, F., & Gehrels, N. 2013, ApJ, 779, 23, doi: 10.1088/0004-637X/779/1/23
  * Mould et al. (2000) Mould, J. R., Huchra, J. P., Freedman, W. L., et al. 2000, ApJ, 529, 786, doi: 10.1086/308304
  * Nagao et al. (2017) Nagao, T., Maeda, K., & Yamanaka, M. 2017, ApJ, 835, 143, doi: 10.3847/1538-4357/835/2/143
  * Niemeyer et al. (1996) Niemeyer, J. C., Hillebrandt, W., & Woosley, S. E. 1996, ApJ, 471, 903, doi: 10.1086/178017
  * Noebauer et al. (2016) Noebauer, U. M., Taubenberger, S., Blinnikov, S., Sorokina, E., & Hillebrandt, W. 2016, MNRAS, 463, 2972, doi: 10.1093/mnras/stw2197
  * Nomoto (1982) Nomoto, K. 1982, ApJ, 253, 798, doi: 10.1086/159682
  * Nomoto et al. (1984) Nomoto, K., Thielemann, F.-K., & Yokoi, K. 1984, ApJ, 286, 644, doi: 10.1086/162639
  * Nugent et al. (1995) Nugent, P., Phillips, M., Baron, E., Branch, D., & Hauschildt, P. 1995, ApJ, 455, L147, doi: 10.1086/309846
  * Oke & Sandage (1968) Oke, J. B., & Sandage, A. 1968, ApJ, 154, 21, doi: 10.1086/149737
  * Oke et al. (1995) Oke, J. B., Cohen, J. G., Carr, M., et al. 1995, PASP, 107, 375, doi: 10.1086/133562
  * Pakmor et al. (2010) Pakmor, R., Kromer, M., Röpke, F. K., et al. 2010, Nature, 463, 61, doi: 10.1038/nature08642
  * Pakmor et al. (2012) Pakmor, R., Kromer, M., Taubenberger, S., et al. 2012, ApJ, 747, L10.  https://arxiv.org/abs/1201.5123
  * Parrent et al. (2011) Parrent, J. T., Thomas, R. C., Fesen, R. A., et al. 2011, ApJ, 732, 30, doi: 10.1088/0004-637X/732/1/30
  * Parrent et al. (2016) Parrent, J. T., Howell, D. A., Fesen, R. A., et al. 2016, MNRAS, 457, 3702, doi: 10.1093/mnras/stw239
  * Patat et al. (2012) Patat, F., Höflich, P., Baade, D., et al. 2012, A&A, 545, A7, doi: 10.1051/0004-6361/201219146
  * Pejcha et al. (2013) Pejcha, O., Antognini, J. M., Shappee, B. J., & Thompson, T. A. 2013, MNRAS, 435, 943, doi: 10.1093/mnras/stt1281
  * Penney & Hoeflich (2014) Penney, R., & Hoeflich, P. 2014, ApJ, 795, 84, doi: 10.1088/0004-637X/795/1/84
  * Perlmutter et al. (1999) Perlmutter, S., Aldering, G., Goldhaber, G., et al. 1999, ApJ, 517, 565, doi: 10.1086/307221
  * Pfannes et al. (2010) Pfannes, J. M. M., Niemeyer, J. C., & Schmidt, W. 2010, A&A, 509, A75, doi: 10.1051/0004-6361/200912033
  * Phillips (1993) Phillips, M. M. 1993, ApJ, 413, L105, doi: 10.1086/186970
  * Phillips et al. (1999) Phillips, M. M., Lira, P., Suntzeff, N. B., et al. 1999, AJ, 118, 1766, doi: 10.1086/301032
  * Phillips et al. (2013) Phillips, M. M., Simon, J. D., Morrell, N., et al. 2013, ApJ, 779, 38, doi: 10.1088/0004-637X/779/1/38
  * Phillips et al. (2019) Phillips, M. M., Contreras, C., Hsiao, E. Y., et al. 2019, PASP, 131, 014001, doi: 10.1088/1538-3873/aae8bd
  * Piascik et al. (2014) Piascik, A. S., Steele, I. A., Bates, S. D., et al. 2014, in Society of Photo-Optical Instrumentation Engineers (SPIE) Conference Series, Vol. 9147, Ground-based and Airborne Instrumentation for Astronomy V, 91478H, doi: 10.1117/12.2055117
  * Piersanti et al. (2003) Piersanti, L., Gagliardi, S., Iben, Icko, J., & Tornambé, A. 2003, ApJ, 598, 1229, doi: 10.1086/378952
  * Poludnenko et al. (2019) Poludnenko, A. Y., Chambers, J., Ahmed, K., Gamezo, V. N., & Taylor, B. D. 2019, Science, 366, aau7365, doi: 10.1126/science.aau7365
  * Poludnenko et al. (2011) Poludnenko, A. Y., Gardiner, T. A., & Oran, E. S. 2011, Physical Review Letters, 107, 054501, doi: 10.1103/PhysRevLett.107.054501
  * Poznanski et al. (2012) Poznanski, D., Prochaska, J. X., & Bloom, J. S. 2012, MNRAS, 426, 1465, doi: 10.1111/j.1365-2966.2012.21796.x
  * Pskovskii (1977) Pskovskii, I. P. 1977, Soviet Ast., 21, 675
  * Quimby et al. (2006) Quimby, R., Höflich, P., Kannappan, S. J., et al. 2006, ApJ, 636, 400, doi: 10.1086/498014
  * Quimby et al. (2007) Quimby, R., Höflich, P., & Wheeler, J. C. 2007, ApJ, 666, 1083, doi: 10.1086/520527
  * Rasmussen & Williams (2005) Rasmussen, C. E., & Williams, C. K. I. 2005, Gaussian Processes for Machine Learning (Adaptive Computation and Machine Learning) (The MIT Press)
  * Riess et al. (1998) Riess, A. G., Filippenko, A. V., Challis, P., et al. 1998, AJ, 116, 1009, doi: 10.1086/300499
  * Riess et al. (2007) Riess, A. G., Strolger, L.-G., Casertano, S., et al. 2007, ApJ, 659, 98, doi: 10.1086/510378
  * Roming et al. (2005) Roming, P. W. A., Kennedy, T. E., Mason, K. O., et al. 2005, Space Sci. Rev., 120, 95, doi: 10.1007/s11214-005-5095-4
  * Rosswog et al. (2009) Rosswog, S., Kasen, D., Guillochon, J., & Ramirez-Ruiz, E. 2009, ApJ, 705, L128, doi: 10.1088/0004-637X/705/2/L128
  * Scalzo et al. (2012) Scalzo, R., Aldering, G., Antilogus, P., et al. 2012, ApJ, 757, 12, doi: 10.1088/0004-637X/757/1/12
  * Scalzo et al. (2010) Scalzo, R. A., Aldering, G., Antilogus, P., et al. 2010, ApJ, 713, 1073, doi: 10.1088/0004-637X/713/2/1073
  * Scalzo et al. (2014) Scalzo, R. A., Childress, M., Tucker, B., et al. 2014, MNRAS, 445, 30, doi: 10.1093/mnras/stu1723
  * Schlafly & Finkbeiner (2011) Schlafly, E. F., & Finkbeiner, D. P. 2011, ApJ, 737, 103, doi: 10.1088/0004-637X/737/2/103
  * Shappee et al. (2014) Shappee, B. J., Prieto, J. L., Grupe, D., et al. 2014, ApJ, 788, 48, doi: 10.1088/0004-637X/788/1/48
  * Sharp & Hoeflich (1990) Sharp, C. M., & Hoeflich, P. 1990, Ap&SS, 171, 213
  * Shen & Moore (2014) Shen, K. J., & Moore, K. 2014, ApJ, 797, 46, doi: 10.1088/0004-637X/797/1/46
  * Silverman et al. (2011) Silverman, J. M., Ganeshalingam, M., Li, W., et al. 2011, MNRAS, 410, 585, doi: 10.1111/j.1365-2966.2010.17474.x
  * Simcoe et al. (2013) Simcoe, R. A., Burgasser, A. J., Schechter, P. L., et al. 2013, PASP, 125, 270, doi: 10.1086/670241
  * Smartt et al. (2015) Smartt, S. J., Valenti, S., Fraser, M., et al. 2015, A&A, 579, A40, doi: 10.1051/0004-6361/201425237
  * Stahl et al. (2020) Stahl, B. E., Zheng, W., de Jaeger, T., et al. 2020, MNRAS, 492, 4325, doi: 10.1093/mnras/staa102
  * Steele et al. (2004) Steele, I. A., Smith, R. J., Rees, P. C., et al. 2004, in Society of Photo-Optical Instrumentation Engineers (SPIE) Conference Series, Vol. 5489, Ground-based Telescopes, ed. J. Oschmann, Jacobus M., 679–692, doi: 10.1117/12.551456
  * Stritzinger et al. (2011) Stritzinger, M. D., Phillips, M. M., Boldt, L. N., et al. 2011, AJ, 142, 156, doi: 10.1088/0004-6256/142/5/156
  * Stritzinger et al. (2015) Stritzinger, M. D., Valenti, S., Hoeflich, P., et al. 2015, A&A, 573, A2, doi: 10.1051/0004-6361/201424168
  * Sullivan et al. (2011) Sullivan, M., Guy, J., Conley, A., et al. 2011, ApJ, 737, 102, doi: 10.1088/0004-637X/737/2/102
  * Suzuki et al. (2012) Suzuki, N., Rubin, D., Lidman, C., et al. 2012, ApJ, 746, 85, doi: 10.1088/0004-637X/746/1/85
  * Tanaka et al. (2010) Tanaka, M., Kawabata, K. S., Yamanaka, M., et al. 2010, ApJ, 714, 1209, doi: 10.1088/0004-637X/714/2/1209
  * Taubenberger (2017) Taubenberger, S. 2017, The Extremes of Thermonuclear Supernovae, ed. A. W. Alsabti & P. Murdin, 317, doi: 10.1007/978-3-319-21846-5_37
  * Taubenberger et al. (2011) Taubenberger, S., Benetti, S., Childress, M., et al. 2011, MNRAS, 412, 2735, doi: 10.1111/j.1365-2966.2010.18107.x
  * Taubenberger et al. (2013) Taubenberger, S., Kromer, M., Hachinger, S., et al. 2013, MNRAS, 432, 3117, doi: 10.1093/mnras/stt668
  * Taubenberger et al. (2019) Taubenberger, S., Floers, A., Vogl, C., et al. 2019, MNRAS, 488, 5473, doi: 10.1093/mnras/stz1977
  * Telesco et al. (2015) Telesco, C. M., Höflich, P., Li, D., et al. 2015, ApJ, 798, 93, doi: 10.1088/0004-637X/798/2/93
  * Thompson (2011) Thompson, T. A. 2011, ApJ, 741, 82, doi: 10.1088/0004-637X/741/2/82
  * Walker et al. (2012) Walker, E. S., Hachinger, S., Mazzali, P. A., et al. 2012, MNRAS, 427, 103, doi: 10.1111/j.1365-2966.2012.21928.x
  * Webbink (1984) Webbink, R. F. 1984, ApJ, 277, 355, doi: 10.1086/161701
  * Wheeler et al. (1998) Wheeler, J. C., Höflich, P., Harkness, R. P., & Spyromilio, J. 1998, ApJ, 496, 908, doi: 10.1086/305427
  * Whelan & Iben (1973) Whelan, J., & Iben, Jr., I. 1973, ApJ, 186, 1007, doi: 10.1086/152565
  * Wyatt et al. (2021) Wyatt, S. D., Sand, D. J., Hsiao, E. Y., et al. 2021, ApJ, 914, 57, doi: 10.3847/1538-4357/abf7c3
  * Yamanaka et al. (2009) Yamanaka, M., Kawabata, K. S., Kinugasa, K., et al. 2009, ApJ, 707, L118, doi: 10.1088/0004-637X/707/2/L118
  * Yamanaka et al. (2016) Yamanaka, M., Maeda, K., Tanaka, M., et al. 2016, PASJ, 68, 68, doi: 10.1093/pasj/psw047
  * Yaron & Gal-Yam (2012) Yaron, O., & Gal-Yam, A. 2012, PASP, 124, 668, doi: 10.1086/666656
  * Yoon & Langer (2005) Yoon, S. C., & Langer, N. 2005, A&A, 435, 967, doi: 10.1051/0004-6361:20042542
  * Yuan et al. (2010) Yuan, F., Quimby, R. M., Wheeler, J. C., et al. 2010, ApJ, 715, 1338, doi: 10.1088/0004-637X/715/2/1338
  * Zhang et al. (2016) Zhang, K., Wang, X., Zhang, J., et al. 2016, ApJ, 820, 67, doi: 10.3847/0004-637X/820/1/67

##  Appendix A Photometry and Observational Logs

Figure A16 presents the $i$-band finding chart of ASASSN-15hy with the
positions of the local-sequence stars, used for CSP-II photometric
observations at LCO. ASAS-SN $V$-band nondetection limits and photometric
observations, in the standard system, are tabulated in Table 4. The CSP-II
optical photometry of ASASSN-15hy is presented in the natural system (see
Section 5.1 of Krisciunas et al. 2017) of the Swope+e2v (Phillips et al.,
2019) in Table 5. The host subtracted $Swift$ UVOT photometry is shown in
Table 6. The CSP-II NIR photometry is presented in the natural system of the
du Pont+RetroCam in Table 7.

Among the previously published optical spectra, twelve were taken with KAST
double spectrograph on the 3-m Shane telescope at the Lick Observatory and the
LRIS (Oke et al., 1995) on the 10-m Keck-I telescope at the W.M. Keck
Observatory and were published by Stahl et al. (2020). Another previously
published series was obtained with the WiFeS instrument (Dopita et al., 2007,
2010) on the Australian National University 2.3-m telescope (ANU) as part of
the ANU WiFeS SuperNovA Programme (AWSNAP, Childress et al. 2016). Note that
red and blue spectra observed on the same night were merged, yielding nine
spectra in total from WiFeS. One classification spectrum taken with the EFOSC2
on the ESO New Technology Telescope (NTT) by the PESSTO program (Smartt et
al., 2015) is also included and is publicly available on WISeREP (Yaron & Gal-
Yam, 2012). A log of the spectroscopic observations of ASASSN-15hy is given in
Table 8. The phases in all tables are expressed in rest-frame phase in days
with respect to $B$-band maximum JD = 2457151.6 $\pm$ 0.4.

Figure A16: The $i$-band finding chart of ASASSN-15hy. The optical (in yellow circles) and NIR (in red circles) local sequence stars are indicated. The location of the SN is labeled with a red line.  Table 4: ASAS-SN $V$-band photometry.aaMagnitudes in the standard system. JD | UT Date | Phase | mag(err)
---|---|---|---
2457100.2 | 2015-03-18 |  $-$50.5 |  $>$17.11
2457108.9 | 2015-03-27 |  $-$42.0 |  $>$17.90
2457115.9 | 2015-04-03 |  $-$35.1 |  $>$16.73
2457116.9 | 2015-04-04 |  $-$34.1 |  $>$16.97
2457117.9 | 2015-04-05 |  $-$33.1 |  $>$17.19
2457129.8 | 2015-04-17 |  $-$21.4 |  $>$18.03
2457133.8 | 2015-04-21 |  $-$17.5 |  $>$17.53
2457137.8 | 2015-04-25 |  $-$13.6 | 16.42(0.10)
2457138.8 | 2015-04-26 |  $-$12.6 | 16.42(0.09)
2457139.8 | 2015-04-27 |  $-$11.6 | 16.12(0.08)
2457143.9 | 2015-05-01 |  $-$7.6 | 15.90(0.06)
2457147.8 | 2015-05-05 |  $-$3.8 | 15.52(0.12)
2457154.8 | 2015-05-12 | 3.1 | 15.37(0.05)
2457155.8 | 2015-05-13 | 4.1 | 15.51(0.05)
2457156.1 | 2015-05-13 | 4.4 | 15.42(0.05)
2457157.1 | 2015-05-14 | 5.4 | 15.42(0.05)
2457158.1 | 2015-05-15 | 6.4 | 15.48(0.05)
2457160.8 | 2015-05-18 | 9.0 | 15.48(0.06)
2457164.7 | 2015-05-22 | 12.8 | 15.51(0.06)
2457168.9 | 2015-05-26 | 17.0 | 15.88(0.06)
2457173.9 | 2015-05-31 | 21.9 | 15.96(0.09)
2457181.8 | 2015-06-08 | 29.6 | 16.30(0.10)
2457182.7 | 2015-06-09 | 30.5 | 16.52(0.14)
2457185.0 | 2015-06-11 | 32.8 | 16.59(0.13)
2457186.8 | 2015-06-13 | 34.5 | 16.45(0.09)
2457189.7 | 2015-06-16 | 37.4 | 16.57(0.13)
2457190.7 | 2015-06-17 | 38.4 | 16.33(0.09)
2457199.1 | 2015-06-25 | 46.6 | 16.72(0.07)
2457228.1 | 2015-07-24 | 75.1 | 17.48(0.13)
2457244.2 | 2015-08-09 | 90.9 | 17.65(0.15)

Table 5: Swope e2v optical photometry of ASASSN-15hy.aaMagnitudes in the Swope+e2v natural system. JD | Phase | $u$ | $g$ | $B$ | $V$ | $r$ | $i$
---|---|---|---|---|---|---|---
2457139.9 |  $-$11.5 | 16.33(0.01) | 16.12(0.01) | 16.21(0.01) | 16.11(0.01) | 16.05(0.01) | 16.17(0.01)
2457140.9 |  $-$10.6 | 16.27(0.01) | 16.03(0.01) | 16.11(0.01) | 16.01(0.01) | 15.96(0.01) | 16.08(0.01)
2457141.9 |  $-$9.6 | 16.22(0.01) | 15.94(0.01) | 16.03(0.01) | 15.93(0.01) | 15.82(0.01) | 15.95(0.01)
2457142.9 |  $-$8.6 | 16.18(0.01) | 15.86(0.01) | 15.96(0.01) | 15.85(0.01) | 15.78(0.02) | …
2457143.8 |  $-$7.6 | 16.18(0.01) | 15.81(0.01) | 15.90(0.01) | 15.78(0.01) | 15.69(0.01) | 15.84(0.01)
2457144.9 |  $-$6.6 | 16.14(0.01) | 15.74(0.01) | 15.83(0.01) | 15.69(0.01) | 15.61(0.01) | 15.79(0.01)
2457145.8 |  $-$5.7 | 16.13(0.01) | 15.70(0.01) | 15.79(0.01) | 15.64(0.01) | 15.59(0.01) | 15.71(0.01)
2457146.9 |  $-$4.7 | 16.14(0.01) | 15.65(0.01) | 15.77(0.01) | 15.60(0.01) | 15.48(0.01) | 15.67(0.01)
2457151.9 | 0.2 | 16.21(0.01) | 15.55(0.01) | 15.69(0.01) | 15.42(0.01) | 15.31(0.01) | 15.55(0.01)
2457152.9 | 1.2 | 16.25(0.01) | 15.53(0.01) | 15.69(0.01) | 15.40(0.01) | 15.29(0.01) | 15.49(0.01)
2457153.9 | 2.2 | 16.28(0.01) | 15.54(0.01) | 15.72(0.01) | 15.40(0.01) | 15.30(0.01) | 15.53(0.01)
2457154.9 | 3.2 | 16.32(0.01) | 15.56(0.01) | 15.75(0.01) | 15.40(0.01) | 15.25(0.01) | 15.47(0.01)
2457155.9 | 4.2 | 16.37(0.01) | 15.56(0.01) | 15.78(0.01) | 15.38(0.01) | 15.27(0.01) | 15.46(0.01)
2457156.9 | 5.1 | 16.40(0.03) | 15.60(0.01) | 15.78(0.02) | 15.40(0.02) | 15.26(0.01) | 15.48(0.01)
2457157.8 | 6.1 | 16.45(0.01) | 15.61(0.01) | 15.84(0.01) | 15.42(0.01) | 15.26(0.01) | 15.46(0.01)
2457158.9 | 7.1 | 16.55(0.01) | 15.65(0.01) | 15.88(0.01) | 15.42(0.01) | 15.27(0.01) | 15.45(0.01)
2457159.9 | 8.1 | 16.64(0.01) | 15.69(0.01) | 15.95(0.01) | 15.44(0.01) | 15.27(0.01) | 15.45(0.01)
2457160.9 | 9.1 | 16.72(0.01) | 15.75(0.01) | 16.00(0.01) | 15.48(0.01) | 15.30(0.01) | 15.47(0.01)
2457161.9 | 10.1 | 16.81(0.01) | 15.79(0.01) | 16.08(0.01) | 15.49(0.01) | 15.27(0.01) | 15.46(0.02)
2457162.9 | 11.0 | 16.91(0.01) | 15.85(0.01) | 16.14(0.01) | 15.52(0.01) | 15.32(0.01) | 15.48(0.01)
2457163.9 | 12.0 | 17.03(0.01) | 15.91(0.01) | 16.20(0.01) | 15.56(0.01) | 15.34(0.01) | 15.46(0.01)
2457164.9 | 13.0 | 17.14(0.01) | 15.97(0.01) | 16.29(0.01) | 15.61(0.01) | 15.35(0.01) | 15.50(0.01)
2457166.9 | 14.9 | 17.35(0.01) | 16.12(0.01) | 16.45(0.01) | 15.70(0.01) | 15.40(0.01) | 15.48(0.01)
2457174.8 | 22.7 | 18.15(0.03) | 16.64(0.01) | 17.05(0.01) | 16.06(0.01) | 15.62(0.01) | 15.62(0.01)
2457175.7 | 23.7 | 18.27(0.04) | 16.70(0.01) | 17.12(0.01) | 16.08(0.01) | 15.64(0.01) | 15.56(0.01)
2457176.8 | 24.7 | 18.25(0.16) | 16.82(0.05) | 17.17(0.06) | 16.14(0.07) | 15.63(0.02) | 15.66(0.04)
2457177.9 | 25.7 | 18.39(0.05) | 16.81(0.01) | 17.22(0.02) | 16.18(0.01) | 15.72(0.01) | 15.61(0.01)
2457178.8 | 26.7 | 18.48(0.03) | 16.88(0.01) | 17.31(0.02) | 16.21(0.01) | 15.74(0.01) | 15.61(0.01)
2457179.8 | 27.7 | 18.58(0.11) | 16.86(0.02) | 17.34(0.04) | 16.23(0.01) | 15.77(0.01) | 15.62(0.01)
2457180.8 | 28.7 | 18.43(0.06) | 16.96(0.01) | 17.36(0.02) | 16.33(0.02) | 15.80(0.02) | 15.67(0.02)
2457181.9 | 29.7 | 18.60(0.03) | 17.00(0.01) | 17.42(0.01) | 16.34(0.01) | 15.84(0.01) | 15.70(0.01)
2457182.9 | 30.7 | 18.65(0.03) | 17.04(0.01) | 17.48(0.01) | 16.37(0.01) | 15.86(0.01) | 15.69(0.01)
2457249.6 | 96.2 | 19.71(0.04) | 18.08(0.01) | 18.48(0.01) | 17.69(0.01) | 17.53(0.01) | …
2457250.6 | 97.2 | 19.65(0.07) | 18.07(0.01) | 18.50(0.02) | 17.71(0.02) | 17.53(0.01) | 17.61(0.02)
Table 6: $Swift$ UVOT photometry of ASASSN-15hy.aaMagnitudes are host subtracted. JD | Phase | $uvw2$ | $uvm2$ | $uvw1$ | $u$ | $b$ | $v$
---|---|---|---|---|---|---|---
2457142.1 |  $-$9.4 | 17.30(0.09) | 17.32(0.09) | 16.33(0.08) | 15.20(0.06) | 16.11(0.06) | 15.98(0.08)
2457145.3 |  $-$6.2 | 17.61(0.10) | 17.35(0.09) | 16.44(0.08) | 15.26(0.06) | 15.88(0.06) | 15.67(0.08)
2457147.3 |  $-$4.3 | … | 17.32(0.16) | 16.32(0.08) | … | … | …
2457147.4 |  $-$4.2 | 17.41(0.10) | 17.29(0.09) | … | 15.21(0.06) | 15.76(0.06) | 15.61(0.07)
2457148.4 |  $-$3.2 | 17.53(0.10) | 17.40(0.09) | 16.47(0.08) | 15.33(0.06) | 15.88(0.06) | 15.62(0.08)
2457149.1 |  $-$2.5 | … | … | … | 15.28(0.06) | … | …
2457149.2 |  $-$2.4 | 17.64(0.10) | 17.51(0.09) | 16.51(0.08) | 15.32(0.06) | 15.70(0.06) | 15.55(0.07)
2457149.5 |  $-$2.1 | 17.73(0.10) | … | 16.54(0.08) | … | … | …
2457149.6 |  $-$2.0 | 17.83(0.11) | 17.36(0.09) | … | 15.34(0.06) | 15.89(0.06) | 15.55(0.07)
2457151.3 |  $-$0.3 | 17.76(0.10) | 17.48(0.10) | 16.50(0.08) | 15.38(0.06) | 15.72(0.06) | 15.57(0.07)
2457154.3 | 2.6 | 17.87(0.12) | 17.91(0.10) | 16.61(0.08) | 15.55(0.07) | 15.88(0.06) | 15.45(0.07)
2457155.3 | 3.6 | 18.08(0.11) | 17.96(0.12) | 16.83(0.08) | 15.55(0.07) | 15.81(0.06) | 15.51(0.07)
2457156.4 | 4.7 | 18.08(0.12) | 17.99(0.11) | 16.89(0.08) | 15.63(0.07) | 15.88(0.06) | 15.45(0.07)
2457157.0 | 5.3 | 18.17(0.12) | 18.22(0.12) | 16.97(0.09) | 15.72(0.06) | 15.95(0.06) | 15.42(0.07)
2457158.4 | 6.6 | 18.23(0.15) | 18.15(0.14) | 16.91(0.10) | 15.73(0.08) | 16.04(0.07) | 15.43(0.08)
2457170.4 | 18.4 | 19.35(0.33) | … | 18.22(0.21) | 17.02(0.14) | 17.04(0.10) | 15.82(0.10)
2457173.0 | 21.0 | … | 20.01(0.36) | 18.76(0.17) | 17.29(0.10) | 17.03(0.07) | 16.04(0.10)
2457175.6 | 23.5 | … | … | 19.05(0.24) | 17.38(0.12) | 17.30(0.09) | 16.13(0.09)
2457181.6 | 29.4 | … | … | 19.19(0.21) | 17.98(0.13) | 17.59(0.08) | 16.39(0.08)
2457181.7 | 29.5 | … | 20.31(0.31) | … | … | … | …
2457183.9 | 31.7 | … | … | 19.39(0.29) | 17.83(0.14) | 17.64(0.10) | 16.48(0.10)
2457185.2 | 33.0 | … | … | 19.32(0.24) | 17.97(0.14) | 17.68(0.09) | …
2457185.3 | 33.1 | … | … | … | … | … | 16.40(0.13)
2457187.8 | 35.5 | … | … | … | 18.26(0.16) | 17.75(0.10) | 16.51(0.09)
2457191.3 | 38.9 | … | … | 19.29(0.23) | 18.27(0.16) | 17.72(0.09) | 16.59(0.09)
2457192.6 | 40.2 | … | … | … | 18.52(0.29) | 18.11(0.16) | 16.97(0.17)
2457197.0 | 44.5 | … | … | … | 18.49(0.19) | 17.77(0.10) | 16.84(0.11)

Table 7: du Pont RetroCam NIR photometry.aaMagnitudes in the du Pont+RetroCam natural system. JD | Phase | $Y$ | $J$ | $H$
---|---|---|---|---
2457139.9 |  $-$11.5 | 15.97(0.01) | 15.82(0.01) | 15.83(0.02)
2457140.9 |  $-$10.5 | 15.88(0.01) | 15.71(0.01) | 15.75(0.02)
2457141.8 |  $-$9.6 | 15.77(0.01) | 15.63(0.01) | 15.61(0.03)
2457142.9 |  $-$8.6 | 15.71(0.01) | 15.54(0.01) | 15.55(0.03)
2457143.9 |  $-$7.6 | 15.65(0.01) | 15.49(0.01) | 15.47(0.03)
2457151.9 | 0.3 | 15.37(0.01) | 15.16(0.01) | 15.19(0.02)
2457152.8 | 1.2 | 15.35(0.01) | 15.17(0.01) | 15.17(0.02)
2457153.8 | 2.2 | 15.35(0.01) | 15.16(0.01) | 15.14(0.02)
2457170.8 | 18.8 | 15.09(0.01) | 15.32(0.01) | 15.03(0.01)
2457172.8 | 20.8 | 15.06(0.01) | 15.34(0.01) | 15.04(0.01)
2457173.8 | 21.8 | 15.03(0.01) | 15.34(0.01) | 15.02(0.02)