# Ground-based and JWST Observations of SN 2022pul:
II. Evidence from Nebular Spectroscopy for a Violent Merger in a Peculiar
Type-Ia Supernova

Lindsey A. Kwok Department of Physics and Astronomy, Rutgers, the State
University of New Jersey,
136 Frelinghuysen Road, Piscataway, NJ 08854-8019, USA  Matthew R. Siebert
Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD
21218-2410, USA  Joel Johansson Oskar Klein Centre, Department of Physics,
Stockholm University, Albanova University Center, SE-106 91, Stockholm, Sweden
Saurabh W. Jha Department of Physics and Astronomy, Rutgers, the State
University of New Jersey,
136 Frelinghuysen Road, Piscataway, NJ 08854-8019, USA  Stéphane Blondin Aix
Marseille Univ, CNRS, CNES, LAM, Marseille, France  Luc Dessart Institut
d’Astrophysique de Paris, CNRS-Sorbonne Université, 98 bis boulevard Arago,
75014, Paris, France  Ryan J. Foley Department of Astronomy and Astrophysics,
University of California, Santa Cruz, CA 95064-1077, USA  D. John Hillier
Department of Physics and Astronomy and Pittsburgh Particle Physics,
Astrophysics and Cosmology Center (PITT PACC), University of Pittsburgh, 3941
O’Hara Street, Pittsburgh, PA 15260, USA  Conor Larison Department of Physics
and Astronomy, Rutgers, the State University of New Jersey,
136 Frelinghuysen Road, Piscataway, NJ 08854-8019, USA  Rüdiger Pakmor Max-
Planck Institute for Astrophysics, Garching, Germany  Tea Temim Department of
Astrophysical Sciences, Princeton University, 4 Ivy Lane, Princeton, NJ
08540-7219, USA  Jennifer E. Andrews Gemini Observatory/NSF’s NOIRLab, 670
North A‘ohoku Place, Hilo, HI 96720-2700, USA  Katie Auchettl School of
Physics, The University of Melbourne, Parkville, VIC 3010, Australia
Department of Astronomy and Astrophysics, University of California, Santa
Cruz, CA 95064-1077, USA  Carles Badenes Department of Physics and Astronomy
and Pittsburgh Particle Physics, Astrophysics and Cosmology Center (PITT
PACC), University of Pittsburgh, 3941 O’Hara Street, Pittsburgh, PA 15260, USA
Barnabas Barna Department of Experimental Physics, Institute of Physics,
University of Szeged, Dóm tér 9, 6720 Szeged, Hungary  K. Azalee Bostroem
Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson,
AZ 85721-0065, USA  Max J. Brenner Newman Department of Physics and Astronomy,
Rutgers, the State University of New Jersey,
136 Frelinghuysen Road, Piscataway, NJ 08854-8019, USA  Thomas G. Brink
Department of Astronomy, University of California, Berkeley, CA 94720-3411,
USA  María José Bustamante-Rosell Department of Astronomy and Astrophysics,
University of California, Santa Cruz, CA 95064-1077, USA  Yssavo Camacho-Neves
Department of Physics and Astronomy, Rutgers, the State University of New
Jersey,
136 Frelinghuysen Road, Piscataway, NJ 08854-8019, USA  Alejandro Clocchiatti
Instituto de Astrofísica, Facultad de Física, Pontificia Universidad Católica
de Chile, Av. Vicuña Mackenna 4860, Santiago, Chile  Millennium Institute of
Astrophysics, Nuncio Monseñor Sótero Sanz 100, Providencia, Santiago, Chile
David A. Coulter Department of Astronomy and Astrophysics, University of
California, Santa Cruz, CA 95064-1077, USA  Kyle W. Davis Department of
Astronomy and Astrophysics, University of California, Santa Cruz, CA
95064-1077, USA  Maxime Deckers School of Physics, Trinity College Dublin, The
University of Dublin, Dublin 2, Ireland  Georgios Dimitriadis School of
Physics, Trinity College Dublin, The University of Dublin, Dublin 2, Ireland
Yize Dong Department of Physics and Astronomy, University of California,
Davis, 1 Shields Avenue, Davis, CA 95616-5270, USA  Joseph Farah Las Cumbres
Observatory, 6740 Cortona Drive, Suite 102, Goleta, CA 93117-5575, USA
Department of Physics, University of California, Santa Barbara, CA 93106-9530,
USA  Alexei V. Filippenko Department of Astronomy, University of California,
Berkeley, CA 94720-3411, USA  Andreas Flörs GSI Helmholtzzentrum für
Schwerionenforschung, Planckstraße 1, 64291 Darmstadt, Germany  Ori D. Fox
Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD
21218-2410, USA  Peter Garnavich Department of Physics and Astronomy,
University of Notre Dame, Notre Dame, IN 46556, USA  Estefania Padilla
Gonzalez Las Cumbres Observatory, 6740 Cortona Drive, Suite 102, Goleta, CA
93117-5575, USA  Department of Physics, University of California, Santa
Barbara, CA 93106-9530, USA  Or Graur Institute of Cosmology & Gravitation,
University of Portsmouth, Dennis Sciama Building, Burnaby Road, Portsmouth PO1
3FX, UK  Department of Astrophysics, American Museum of Natural History,
Central Park West and 79th Street, New York, NY 10024-5192, USA  Franz-Josef
Hambsch Vereniging Voor Sterrenkunde (VVS), Oostmeers 122 C, 8000 Brugge,
Belgium  AAVSO, 185 Alewife Brook Parkway, Suite 410, Cambridge, MA 02138, USA
Groupe Européen d’Observations Stellaires (GEOS), 23 Parc de Levesville, 28300
Bailleau l’Evêque, France  Bundesdeutsche Arbeitsgemeinschaft für
Veränderliche Sterne (BAV), Munsterdamm 90, 12169 Berlin, Germany  Griffin
Hosseinzadeh Steward Observatory, University of Arizona, 933 North Cherry
Avenue, Tucson, AZ 85721-0065, USA  D. Andrew Howell Las Cumbres Observatory,
6740 Cortona Drive, Suite 102, Goleta, CA 93117-5575, USA  Department of
Physics, University of California, Santa Barbara, CA 93106-9530, USA  John P.
Hughes Department of Physics and Astronomy, Rutgers, the State University of
New Jersey,
136 Frelinghuysen Road, Piscataway, NJ 08854-8019, USA  Wolfgang E. Kerzendorf
Department of Physics and Astronomy, Michigan State University, East Lansing,
MI 48824, USA  Department of Computational Mathematics, Science, and
Engineering, Michigan State University, East Lansing, MI 48824, USA  Xavier K.
Le Saux Department of Astronomy and Astrophysics, University of California,
Santa Cruz, CA 95064-1077, USA  Keiichi Maeda Department of Astronomy, Kyoto
University, Kitashirakawa-Oiwake-cho, Sakyo-ku, Kyoto, 606-8502. Japan  Kate
Maguire School of Physics, Trinity College Dublin, The University of Dublin,
Dublin 2, Ireland  Curtis McCully Las Cumbres Observatory, 6740 Cortona Drive,
Suite 102, Goleta, CA 93117-5575, USA  Department of Physics, University of
California, Santa Barbara, CA 93106-9530, USA  Cassidy Mihalenko School of
Physics, The University of Melbourne, Parkville, VIC 3010, Australia  ARC
Centre of Excellence for All Sky Astrophysics in 3 Dimensions (ASTRO 3D)
Megan Newsome Las Cumbres Observatory, 6740 Cortona Drive, Suite 102, Goleta,
CA 93117-5575, USA  Department of Physics, University of California, Santa
Barbara, CA 93106-9530, USA  John T. O’Brien Department of Physics and
Astronomy, Michigan State University, East Lansing, MI 48824, USA  Jeniveve
Pearson Steward Observatory, University of Arizona, 933 North Cherry Avenue,
Tucson, AZ 85721-0065, USA  Craig Pellegrino Las Cumbres Observatory, 6740
Cortona Drive, Suite 102, Goleta, CA 93117-5575, USA  Department of Physics,
University of California, Santa Barbara, CA 93106-9530, USA  Justin D. R.
Pierel Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD
21218-2410, USA  Abigail Polin Observatories of the Carnegie Institute for
Science, 813 Santa Barbara Street, Pasadena, CA 91101-1232, USA  TAPIR, Walter
Burke Institute for Theoretical Physics, 350-17, Caltech, Pasadena, CA 91125,
USA  Armin Rest Space Telescope Science Institute, 3700 San Martin Drive,
Baltimore, MD 21218-2410, USA  Department of Physics and Astronomy, The Johns
Hopkins University, 3400 North Charles Street, Baltimore, MD 21218, USA  César
Rojas-Bravo Department of Astronomy and Astrophysics, University of
California, Santa Cruz, CA 95064-1077, USA  David J. Sand Steward Observatory,
University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA
Michaela Schwab Department of Physics and Astronomy, Rutgers, the State
University of New Jersey,
136 Frelinghuysen Road, Piscataway, NJ 08854-8019, USA  Melissa Shahbandeh
Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD
21218-2410, USA  Manisha Shrestha Steward Observatory, University of Arizona,
933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  Nathan Smith Steward
Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ
85721-0065, USA  Louis-Gregory Strolger Space Telescope Science Institute,
3700 San Martin Drive, Baltimore, MD 21218-2410, USA  Tamás Szalai Department
of Experimental Physics, Institute of Physics, University of Szeged, Dóm tér
9, 6720 Szeged, Hungary  ELKH-SZTE Stellar Astrophysics Research Group,
Szegedi út, Kt. 766, 6500 Baja, Hungary  Kirsty Taggart Department of
Astronomy and Astrophysics, University of California, Santa Cruz, CA
95064-1077, USA  Giacomo Terreran Las Cumbres Observatory, 6740 Cortona Drive,
Suite 102, Goleta, CA 93117-5575, USA  Department of Physics, University of
California, Santa Barbara, CA 93106-9530, USA  Jacco H. Terwel School of
Physics, Trinity College Dublin, The University of Dublin, Dublin 2, Ireland
Isaac Newton Group (ING), Apt. de correos 321, E-38700, Santa Cruz de La
Palma, Canary Islands, Spain  Samaporn Tinyanont National Astronomical
Research Institute of Thailand, 260 Moo 4, Donkaew, Maerim, Chiang Mai 50180,
Thailand  Stefano Valenti Department of Physics and Astronomy, University of
California, Davis, 1 Shields Avenue, Davis, CA 95616-5270, USA  József Vinkó
Konkoly Observatory, Research Centre for Astronomy and Earth Sciences (CSFK),
MTA Center of Excellence, Konkoly-Thege Miklós út 15-17, 1121 Budapest,
Hungary  Department of Experimental Physics, Institute of Physics, University
of Szeged, Dóm tér 9, 6720 Szeged, Hungary  ELTE Eötvös Loránd University,
Institute of Physics and Astronomy, Pázmány Péter sétány 1/A, Budapest, 1117
Hungary  Department of Astronomy, University of Texas at Austin, Austin, TX,
USA  J. Craig Wheeler Department of Astronomy, University of Texas at Austin,
Austin, TX, USA  Yi Yang Department of Astronomy, University of California,
Berkeley, CA 94720-3411, USA  WeiKang Zheng Department of Astronomy,
University of California, Berkeley, CA 94720-3411, USA  Chris Ashall
Department of Physics, Virginia Tech, Blacksburg, VA 24061, USA  James M.
DerKacy Department of Physics, Virginia Tech, Blacksburg, VA 24061, USA  Lluís
Galbany Institute of Space Sciences (ICE, CSIC), Campus UAB, Carrer de Can
Magrans, s/n, E-08193 Barcelona, Spain  Institut d’Estudis Espacials de
Catalunya (IEEC), E-08034 Barcelona, Spain  Peter Hoeflich Department of
Physics, Florida State University, 77 Chieftan Way, Tallahassee, FL
32306-4350, USA  Thomas de Jaeger LPNHE, (CNRS/IN2P3, Sorbonne Université,
Université Paris Cité), Laboratoire de Physique Nucléaire et de Hautes
Énergies, 75005, Paris, France  Jing Lu Department of Physics and Astronomy,
Michigan State University, East Lansing, MI 48824, USA  Justyn Maund
Department of Physics and Astronomy, University of Sheffield, Hicks Building,
Hounsfield Road, Sheffield S3 7RH, U.K.  Kyle Medler Astrophysics Research
Institute, Liverpool John Moores University, Liverpool, L3 5RF, UK  Nidia
Morrell Carnegie Observatories, Las Campanas Observatory, Casilla 601, La
Serena, Chile  Benjamin J. Shappee Institute for Astronomy, University of
Hawai‘i, 2680 Woodlawn Drive, Honolulu, HI 96822-1839, USA  Maximilian
Stritzinger Department of Physics and Astronomy, Aarhus University, Ny
Munkegade 120, DK-8000 Aarhus C, Denmark  Nicholas Suntzeff Department of
Physics and Astronomy, Texas A&M University, 4242 TAMU, College Station, TX
77843, USA  George P. and Cynthia Woods Mitchell Institute for Fundamental
Physics & Astronomy, College Station, TX 77843, USA  Michael Tucker Center for
Cosmology and Astroparticle Physics, The Ohio State University, 191 West
Woodruff Ave, Columbus, OH 43215, USA  Lifan Wang Department of Physics and
Astronomy, Texas A&M University, 4242 TAMU, College Station, TX 77843, USA
George P. and Cynthia Woods Mitchell Institute for Fundamental Physics &
Astronomy, College Station, TX 77843, USA

###### Abstract

We present an analysis of ground-based and JWST observations of SN 2022pul, a
peculiar “03fg-like” (or “super-Chandrasekhar”) Type Ia supernova (SN Ia), in
the nebular phase at 338 d post explosion. Our combined spectrum continuously
covers 0.4–14 $\mu$m and includes the first mid-infrared spectrum of an 03fg-
like SN Ia. Compared to normal SN Ia 2021aefx, SN 2022pul exhibits a lower
mean ionization state, asymmetric emission-line profiles, stronger emission
from the intermediate-mass elements (IMEs) argon and calcium, weaker emission
from iron-group elements (IGEs), and the first unambiguous detection of neon
in a SN Ia. Strong, broad, centrally peaked [Ne II] at 12.81 $\mu$m was
previously predicted as a hallmark of “violent merger” SN Ia models, where
dynamical interaction between two sub-$M_{\text{Ch}}$ white dwarfs (WDs)
causes disruption of the lower mass WD and detonation of the other. The
violent merger scenario was already a leading hypothesis for 03fg-like SNe Ia;
in SN 2022pul it can explain the large-scale ejecta asymmetries seen between
the IMEs and IGEs and the central location of narrow oxygen and broad neon. We
modify extant models to add clumping of the central ejecta to better reproduce
the optical iron emission, and add mass in the innermost region ($<2000$ km
s-1) to account for the observed narrow [O I] $\lambda\lambda 6300$, 6364
emission. A violent WD-WD merger explains many of the observations of SN
2022pul, and our results favor this model interpretation for the subclass of
03fg-like SN Ia.

Supernovae (1668), Type Ia supernovae (1728), White dwarf stars (1799)

††journal: The Astrophysical Journal††facilities: AAVSO, ANU (WiFeS), ASAS-SN,
ATLAS, GTC (OSIRIS), JWST (NIRSpec/MIRI), Keck:I (LRIS), Keck:II (NIRES),
Keck:II (DEIMOS), LCO/GSP, Magellan (IMACS), MMT (Binospec), SALT (RSS), Shane
(Kast), SOAR (Goodman), Subaru (FOCAS), UH (SNIFS), ZTF††software: Astropy
(Astropy Collaboration et al., 2013, 2018, 2022), Matplotlib (Hunter, 2007),
NumPy (Oliphant, 2006), PyRAF (Science Software Branch at STScI, 2012), PySALT
(Crawford et al., 2010), dust extinction (Gordon et al., 2022), jdaviz (JDADF
Developers et al., 2022), jwst (Bushouse et al., 2022), UltraNest (Buchner,
2021; Buchner et al., 2022), YSE-PZ (Coulter et al., 2022, 2023), CMFGEN
(Hillier & Dessart, 2012) ††thanks: LSSTC Catalyst Fellow††thanks: CCAPP
Fellow

##  1 Introduction

Supernovae (SNe) imprint important information about the composition,
geometry, kinematics, and ionization of the ejected material in their late-
time (nebular) emission-line spectra. In the nebular phase ($\gtrsim 100$ days
past maximum light; Bowers et al., 1997; Branch et al., 2008; Silverman et
al., 2013; Friesen et al., 2014; Black et al., 2016), the ejecta have expanded
and the opacity has dropped enough to see directly through even the innermost
regions and allow forbidden lines to form. In white-dwarf (WD) supernovae,
decay of radioactive 56Co (produced by the radioactive decay of 56Ni
synthesized in the explosion) provides most of the energy that excites the
low-density nebular ejecta, producing forbidden emission lines throughout the
ejecta (Axelrod, 1980; Nadyozhin, 1994; Jerkstrand, 2017). These glowing ashes
contain clues for determining the SN origin. Nebular-phase spectroscopy is
therefore an important forensic tool for investigating the remains of exploded
WDs, whose explosion channels are not well understood.

The majority of WD SNe are classified as Type Ia supernovae (SNe Ia); however,
a growing number and variety of extreme or peculiar SN Ia subtypes have been
discovered (for recent reviews, see Taubenberger, 2017; Jha et al., 2019; Liu
et al., 2023). These peculiar SNe result from the thermonuclear explosion of a
WD, but they generally do not follow the Phillips (1993) relationship between
luminosity and light-curve shape that makes normal SNe Ia so useful for
cosmological distance measurements (Phillips, 1993; Benetti et al., 2005;
Taubenberger, 2017; Jha et al., 2019). No consensus has been reached
concerning the possible and dominant pathways connecting progenitor systems,
explosion mechanisms, and observed WD SN properties.

The origins of normal SNe Ia are particularly elusive in part because
variations between objects are relatively small and many of the models predict
similar observational properties. By contrast, extreme and peculiar WD SNe
exhibit distinguishing features that can make identification of their origins
clearer (Foley et al., 2013; Kromer et al., 2013; Hsiao et al., 2020).
Understanding the scenarios that produce peculiar WD SNe helps us understand
the unusual and extreme ways WDs can explode while simultaneously narrowing
the model space for normal SNe Ia.

A rare type of peculiar WD SN characterized by high luminosities, broad light
curves, C II absorption at early times, relatively low ejecta velocities, and
low ionization state nebular-phase spectra are “03fg-like” (commonly called
“super-Chandrasekhar”) SNe Ia (Howell et al., 2006; Hicken et al., 2007;
Scalzo et al., 2010; Silverman et al., 2011; Taubenberger et al., 2013; Hsiao
et al., 2020; Ashall et al., 2021; Srivastav et al., 2023). Explaining the
light curve of SN 2003fg, the class prototype, by decay of 56Ni alone requires
a mass above 1.4  $M_{\odot}$, the Chandrasekhar mass limit ($M_{\text{Ch}}$)
for a nonrotating WD (Howell et al., 2006). A mass above this limit can be
achieved either by additional support from differential rotation in a single
WD (Yoon & Langer, 2005) or a merger of two WDs with total mass greater than
$M_{\text{Ch}}$ (Howell et al., 2006; Hicken et al., 2007). Alternative
explanations that have been invoked to explain this diverse class of bright
SNe Ia include the “core-degenerate” scenario, where a near-$M_{\text{Ch}}$
C/O WD explodes within the C-rich envelope of an asymptotic giant branch (AGB)
star (Khokhlov et al., 1993; Hoeflich et al., 1996; Livio & Riess, 2003;
Quimby et al., 2007; Kashi & Soker, 2011; Hsiao et al., 2020; Lu et al.,
2021), and the violent merger of two sub-$M_{\text{Ch}}$ WDs (Dimitriadis et
al., 2022; Srivastav et al., 2023; Siebert et al., 2023a). These potential
avenues account for the excess luminosity and relatively low velocity of these
objects by interaction with the AGB star envelope or C/O-rich circumstellar
material (CSM) from the disrupted smaller WD, respectively. For more
comprehensive reviews of 03fg-like SNe Ia, see Taubenberger (2017), Ashall et
al. (2021), and Paper I (Siebert et al., 2023b).

In contrast to the classical WD merger scenario where the merger produces a
single-object remnant that later explodes, the violent merger model proceeds
via dynamical interaction between two sub-$M_{\text{Ch}}$ WDs that causes
disruption of the secondary (lower mass) WD and subsequent detonation of the
primary (higher mass) WD (for a review, see Pakmor, 2017). The primary WD is
burned first and its ejecta expand beyond the secondary by the time that the
secondary is burned $\sim 1$ s later and the ejecta of the secondary expand
into the center of the ejecta of the primary. Thus, the intermediate-mass
elements (IMEs) produced by burning of the secondary will be centrally located
within the ejecta. Depending on the mass ratio, mass of the primary WD,
density of the secondary WD at the time of detonation, and composition of both
WDs, a fairly wide range of SN properties can be produced (Pakmor et al.,
2011, 2012; Kromer et al., 2013; Pakmor, 2017). In general, violent mergers
are asymmetric scenarios which can result in enhanced production of IMEs from
burning of the secondary, lower levels of iron-group elements (IGEs) from
burning of the sub-$M_{\text{Ch}}$ primary, centrally located C/O/Ne/Mg and
IMEs, and a low ionization state at late times due to high central ejecta
density.

Nebular spectra can reveal all of these properties. This information can be
difficult to disentangle, however, owing to significant line crowding and
blending at optical wavelengths. Spectroscopy with JWST in the near-infrared
(NIR, which extends beyond the NIR wavelength range accessible from the
ground) and mid-infrared (MIR) grants access to emission lines from several
species of IMEs, and the lines are comparatively isolated so individual line
profiles can be distinguished. These NIR and MIR spectral features help
clarify the information contained in the optical and ground-based NIR spectra,
aiding in discriminating between progenitor systems and explosion mechanisms.
Observations of SN 2021aefx, the first SN Ia with JWST NIR$+$MIR nebular
spectra, revealed strong, stable 58Ni indicative of a high-mass WD progenitor
and evidence for stratified ejecta resulting from some type of detonation
(Kwok et al., 2023). Further modeling found SN 2021aefx to be consistent with
an off-center delayed detonation of a near-$M_{\text{Ch}}$ mass WD (DerKacy et
al., 2023).

Figure 1: Full optical$+$NIR$+$MIR comparison between the dust-continuum-
subtracted spectrum of SN 2022pul at 338 rest-frame days post-explosion
($t_{\text{exp}}=$ MJD 59785.3; $d=16$ Mpc), and SN 2021aefx at 270 d rest-
frame days post-explosion ($d=18$ Mpc) from Kwok et al. (2023). The optical
spectrum of SN 2022pul is from MMT/Binospec at a similar epoch of 332 d post-
explosion and the NIR$+$MIR spectrum is from JWST/NIRSpec$+$MIRI. The
prominent stable 58Ni emission lines in SN 2021aefx are weaker in SN 2022pul,
as is the [Co III] 11.88 $\mu$m line. In SN 2022pul, the singly ionized
emission lines are relatively strong compared with the higher ionization
states. A strong [Ne II] feature at 12.81 $\mu$m is observed in SN 2022pul.
The flux axis uses a non-linear (arcsinh) scale to better show all the
features across a wide range of wavelength and $F_{\nu}$.

We obtained a JWST NIR$+$MIR nebular spectrum of SN 2022pul, a peculiar 03fg-
like SN Ia (JWST Cycle 1 GO 2072; PI S. W. Jha). Our JWST spectrum of SN
2022pul is the first MIR spectrum of an 03fg-like SN Ia, and it exhibits a
strong thermal ($T\approx 500$ K) dust continuum (analyzed in Paper III;
Johansson et al., 2023). It also displays several differences from the
optical–MIR spectrum of SN 2021aefx. Additional details about SN 2022pul and
its unique photometric and spectroscopic evolution can be found in Paper I
(Siebert et al., 2023b).

Here, we present analysis of the dust-continuum-subtracted optical through MIR
nebular-phase spectrum of SN 2022pul at 338 d post explosion (this assumes
that the dust is external, see Johansson et al. 2023 for a discussion). In
Section 2 we identify NIR and MIR emission lines and highlight distinctive
spectral properties in comparison to SN 2021aefx. Section 3 presents emission-
line-profile fits to the dominant spectral features and their implications on
the distribution of material in the ejecta. In Section 4 we compare SN 2022pul
to nebular spectroscopic predictions (Blondin et al., 2023) of the violent
merger model from Pakmor et al. (2012). We discuss the implications of our
results and conclude in Section 5 that SN 2022pul was most likely produced by
the violent merger of two WDs.

##  2 Spectral Analysis

The continuous optical$+$NIR$+$MIR spectrum of SN 2022pul in the nebular phase
at 338 rest-frame days post-explosion (estimated explosion date MJD 59785.3)
exhibits many unique properties, including a clear NIR$+$MIR dust continuum
that closely resembles a blackbody (see also Siebert et al., 2023b). Here, we
focus on analysis of the spectral lines, while a detailed analysis of the dust
properties is given by Johansson et al. (2023). We remove the dust
contribution to the spectrum by subtracting off a blackbody of $T=500$ K from
the 2.5–14 $\mu$m region, and use this dust-subtracted spectrum throughout
this work. We do not observe other spectral signs of dust such as CO or SiO
emission (although line overlap may complicate clear identification of dust
extinction in the optical). For details of the observations and data reduction
of the spectra used in this work, see Siebert et al. (2023b).

Shown in Figure 1, several highlights of SN 2022pul’s combined nebular
spectrum include a low mean ionization state, strong emission from IMEs,
relatively weak emission from IGEs, highly asymmetric line profiles, narrow [O
I] $\lambda\lambda 6300,~{}6364$ and remarkably strong [Ca II] $\lambda\lambda
7291,~{}7324$ in the optical spectrum, detection of [S IV] at 10.51 $\mu$m,
and the first unambiguous detection of [Ne II] at 12.81 $\mu$m in an SN Ia.

Figure 2: Comparison between the dust-continuum-subtracted MIR spectrum of SN
2022pul at 338 rest-frame days post-explosion (MJD 59785.3) and the MIR
spectrum of SN 2021aefx at 340d rest-frame days post-explosion from DerKacy et
al. (2023) scaled to the distance of SN 2022pul (16 Mpc). The prominent stable
58Ni and Co emission lines in SN 2021aefx are somewhat weaker and have shifted
to lower mean ionization state in SN 2022pul. The [Ar II] 6.98 $\mu$m line
dominates the SN 2022pul spectrum, with strong [Ar III] 8.99 $\mu$m and [S IV]
10.51 $\mu$m as well. [Ne II] at 12.81 $\mu$m is strong, broad, and centrally
peaked in SN 2022pul. All line profiles are considerably asymmetric. An inset
is provided on a log scale to show the very weak [Ni IV] 8.41 $\mu$m and [Ni
III] 11.00 $\mu$m lines.

###  2.1 Spectroscopic Comparison with SN 2021aefx

SN 2021aefx, a spectroscopically and photometrically regular SN Ia near
maximum light, is currently the only other SN Ia with a published continuous
spectrum from 0.3 $\mu$m to 14 $\mu$m in the nebular phase (270 d post-
explosion Kwok et al., 2023), making it a natural comparison object for the
peculiar “super-Chandra” 03fg-like SN 2022pul. After subtracting the dust
continuum, SN 2022pul has a greater fraction of total flux in the NIR and MIR,
and a smaller fraction of optical flux than SN 2021aefx. The light curve of SN
2022pul shows that at late times, the I-band flux exceeds the flux in other
bands (probably due to the strong [Ca II] $\lambda\lambda 7291,~{}7324$
doublet) and declines more slowly than typical SNe Ia (Siebert et al., 2023b).
This is consistent with findings from Ashall et al. (2021) that 03fg-like SNe
Ia are brighter and decline more slowly in the NIR than normal SNe Ia.

This redder color may be due in part to the comparatively low overall
ionization state of SN 2022pul, where additional flux from lines of lower
ionization falls in the infrared. The flux from optical [Fe III] around 5000 Å
is greatly diminished relative to most SNe Ia (Siebert et al., 2023b);
similarly, the JWST NIRSpec spectrum reveals that the NIR [Fe III] lines
around 2.2 and 2.9 $\mu$m are weak while the NIR [Fe II] lines around 1.3 and
1.6 $\mu$m are strong compared with SN 2021aefx. In the MIR, we see the same
ionization trend in the Ni, Ar, and Co: the ratio in peak flux between the
singly ionized and doubly ionized lines, $n^{+}/n^{2+}$, is higher in SN
2022pul than in SN 2021aefx. In agreement with the ionization inferred from
the MIR, the strength of the [Ca II] $\lambda\lambda 7291,~{}7324$ doublet in
SN 2022pul is remarkable (Siebert et al., 2023b), whereas in SN 2021aefx the
spectral complex near 7300 Å was attributed to [Fe II] and [Ni II] with little
to no contribution from [Ca II] (Kwok et al., 2023) — yet they have similar
strengths of [Ca IV] near 3.2 $\mu$m. This shift to a lower mean ionization
for SN 2022pul suggests a larger density or mass of the ejecta, but the
presence of Ca IV also indicates that some higher ionization bubbles are
present in both, discussed further in Section 4.

The NIR does not display lines other than those present in SN 2021aefx, though
the [Fe III] is weaker and the three [Fe II] features in the range 4–5 $\mu$m
are stronger. Blondin et al. (2023) showed that the broad feature near 3.2
$\mu$m is dominated by [Ca IV] 3.21 $\mu$m, with weaker contribution from [Fe
III] 3.23 $\mu$m and [Ni I] 3.12 $\mu$m. The NIRSpec observation of SN 2022pul
further confirms this broad line around 3.2 $\mu$m as [Ca IV] 3.21 $\mu$m
because the [Fe III] and Ni emission are weaker in SN 2022pul than in SN
2021aefx, but the 3.2 $\mu$m feature has a similar strength. Several unique
features in the MIRI/LRS spectrum of SN 2022pul are shown in Figure 2 and
described below.

###  2.2 [Ne II] emission

In addition to all of the NIR and MIR lines present in SN 2021aefx (line
identifications are given by Kwok et al. 2023), SN 2022pul exhibits a strong,
broad, centrally peaked feature around 12.8 $\mu$m that we identify as [Ne II]
12.81 $\mu$m (see Figure 2). Predicted by Blondin et al. (2023), the presence
of strong, centrally peaked [Ne II] 12.81 $\mu$m is a distinguishing feature
of a violent WD–WD merger not present in any of the other models considered in
their study. Ne is produced centrally in a violent WD–WD merger because it is
a product of burning the less dense secondary WD, whose ejecta expands into
the center of the primary WD’s ejecta after the primary has burned and
expanded. The observed [Ne II] 12.81 $\mu$m line is broad, asymmetric, and
dissimilar to other line-profile shapes seen in SN 2022pul and other SNe Ia
with MIR spectra.

###  2.3 Strong Intermediate-Mass Elements

One of the major differences between the MIR spectra of SN 2022pul and SN
2021aefx is the strength of the IME (hereafter defined as S, Ar, and Ca)
emission lines. Whereas in SN 2021aefx, [Ar II] 6.98 $\mu$m was blended out by
the neighboring Ni emission lines such that its shape was uncertain, [Ar II]
6.98 $\mu$m is the dominant emission line in the MIR spectrum of SN 2022pul.
Its peak flux is blueshifted by $\sim 2000$ km s-1, and is over twice as
strong in $F_{\nu}$ as [Ne II] 12.81 $\mu$m, the next strongest MIR emission
line. Similarly, [Ca II] $\lambda\lambda 7291,~{}7324$ dominates the optical
spectrum.

The [Ar III] 8.99 $\mu$m emission line is also quite strong in SN 2022pul:
roughly three times stronger in $F_{\nu}$ than in SN 2021aefx at a very
similar time post-explosion (see Figure 2). The [Ar III] 8.99 $\mu$m line
continues to exhibit a boxy, flat-topped shape, with a similar width, but the
slanted top is much steeper. This steep slant can be produced by an asymmetric
shell of [Ar III] emission, discussed further in Section 3.

The feature at 10.5 $\mu$m in SN 2022pul results from a combination of [S IV]
10.51 $\mu$m and [Co II] 10.52 $\mu$m emission. The strength and flat shape of
the feature at 10.5 $\mu$m cannot be explained by [Co II] 10.52 $\mu$m alone,
as is further discussed in Section 3, constituting a firm detection of [S IV].
Kwok et al. (2023) did not invoke a contribution from [S IV] 10.51 $\mu$m in
SN 2021aefx; however, placed in context with SN 2022pul, the “shoulder”
feature starting at 10 $\mu$m in SN 2021aefx might actually be [Si IV] with a
similar boxy shape as the [Ar III] but with about half the strength. Re-
evaluation of [S IV] in SN 2021aefx will be included in future work.

All of the IME emission lines in SN 2022pul are strong relative to SN
2021aefx, and their peak emission is asymmetric and blueshifted.

###  2.4 Weak Iron-group Elements

In contrast to the strength of the IMEs, the IGEs (hereafter Fe, Co, and Ni)
appear somewhat weaker in the MIR than in SN 2021aefx. Part or all of this
apparent decrease in IGEs may be due to the shift to lower ionization. For
example, the [Ni II] 6.64 $\mu$m line is much stronger in SN 2022pul, but the
[Ni III] lines at 7.35 and 11.00 $\mu$m are several times weaker, and the [Ni
IV] 8.41 $\mu$m line is so weak that it is difficult to see without an arcsinh
or log scaling of the flux. This points toward some of the flux in the higher
ionization IGE emission lines being shifted to the singly ionized emission
lines.

The isolated [Co III] 11.89 $\mu$m line shows a distinctly asymmetric shape
whose peak is redshifted by $\sim 3500$ km s-1. Compared with SN 2021aefx, the
[Co III] 11.89 $\mu$m line is somewhat narrower. As will be shown and
discussed in Section 3, the asymmetric redshifted shape of the [Co III] 11.89
$\mu$m line profile is consistent with all of the other IGE ions throughout
the NIR and MIR, as well as the redder side of the optical.

##  3 Emission-Line Profiles

Figure 3: [Ar II] and [Ca II] emission-line profiles in velocity space
compared with [Co III] and [Fe II]. The shapes of the [Ar II] and [Ca II] are
quite consistent with each other, and the shapes of [Co III] and [Fe II] are
nearly identical (taking into account the neighboring [Fe II] lines that
create the red tail on the [Fe II] 1.26 $\mu$m feature). Notably, the Ar and
Ca (IMEs) have a blueshifted peak with a more extended red tail, while the Co
and Fe (IGEs) display an opposite trend, with a redshifted peak and a more
extended blue tail.

Nebular emission-line profiles are important indicators of the geometry of the
SN ejecta because they imprint the flux at all projected line-of-sight
velocities (and therefore projected radius, given homologous expansion)
through the ejecta. The shape of the lines arises from a combination of the
density and excitation of the elements within the ejecta (for a review, see
Jerkstrand, 2017). In SN 2022pul, we find evidence in the highly asymmetric
line profiles for distinct distributions of the IGEs (Fe, Co, Ni), IMEs (S,
Ar, Ca), O, and Ne within the ejecta. Unfortunately, it is not possible to
determine the full three-dimensional (3D) geometry that gives rise to each
profile owing to projection effects and the possibility of bubbles, mushroom
shapes, clumps, or other such non-spherically symmetric distributions.
However, by fitting line profiles throughout the optical, NIR, and MIR, we can
determine the composition within slabs perpendicular to the line of sight,
which resolves some of the geometrical layout. In this section we fit line
profiles consistently across the optical+NIR+MIR and discuss their
implications.

For our line-profile fitting, we need relatively high resolution to discern
line-profile shapes, so we combine an optical MMT/Binospec spectrum at 332 d
post-explosion, a NIR Keck/NIRES spectrum at 316 d post-explosion (the
nearest-phase ground-based NIR spectrum available), and the
JWST/NIRSpec$+$MIRI spectrum beyond 1.83 $\mu$m at 338 d post-explosion
(Siebert et al., 2023b). Despite the difference in phase, we use the
Keck/NIRES spectrum (rebinned to increase the signal-to-noise ratio per pixel)
instead of the JWST/NIRSpec spectrum in the range 1.0–1.83 $\mu$m owing to the
much higher resolution. In addition to removing the dust contribution, we also
remove all remaining underlying MIR continuum in the JWST data (see Johansson
et al., 2023, for discussion of the MIR continuum).

The optical, NIR, and MIR spectra of SN 2022pul are all essential to this
work: the higher resolution of the ground-based optical and NIR spectra
provides clarity for several key line-profile shapes, and the lower resolution
JWST data contain isolated lines and emission from IMEs. These properties of
the data are complementary for interpreting the line profiles of SN 2022pul.

###  3.1 Line-Fitting Procedure

Following the general approach of nebular line fitting from Maguire et al.
(2018), Flörs et al. (2018, 2020), and Kwok et al. (2023), we model the
observed spectrum as a superposition of line profiles for all reasonably
strong contributing lines. Restrictions on the fitting are enforced such that
all lines of the same ion have the same profile shape parameters and kinematic
offset from the central wavelength. We note similarities between the nebular
spectrum of SN 2022pul and the nebular spectrum of the violent WD–WD merger
model of Pakmor et al. (2012) computed by Blondin et al. (2023) (their MERGER
model). Thus, we compute this model out to the same phase as SN 2022pul (338 d
post-explosion) and we fix the relative line strengths of each contributing
line in our fits to the relative line strengths in the model. This allows us
to base the relative line strengths on temperatures and densities similar to
those in SN 2022pul. Details of and comparisons to the merger model are
presented in Section 4.

While the relative line strengths within a given ion are fixed by the MERGER
model (Blondin et al., 2023) at 338 d, we allow the relative amplitudes
between each ion to vary as fitting parameters. We fit the optical through MIR
and include the following ions: Fe I, Fe II, Fe III, Co II, Co III, Ni II, Ni
III, Ni IV, S IV, Ca II, Ca IV, Ar II, Ar III, O I, and Ne II. Additionally,
to account for asymmetry in SN 2022pul, we fit each ion with non-Gaussian line
profiles, further described below.

The resolution of the NIRSpec/Prism and MIRI/LRS data introduces significant
uncertainties into the fits to features beyond 1.8 $\mu$m, on the order of
$1200$ km s-1 near 3 $\mu$m, $2000$ km s-1 near 6 $\mu$m, and $500$ km s-1
near 12 $\mu$m. Uncertainties from the fitting itself are $\sim 1000$ km s-1
for the kinematic offsets and $\sim 1500$ km s-1 for the widths, with the more
isolated lines having somewhat lower uncertainties and the more blended
features having higher uncertainties.

Figure 4: Comparison of the [Ar II] 6.98 $\mu$m and [Ar III] 8.99 $\mu$m
lines, as well as the optical [Ca II] $\lambda\lambda 7291,~{}7324$ and NIR
[Ca IV] 3.21 $\mu$m lines. The lower ionization states show a more prominent
blueshifted peak (vertical dashed blue line), while the higher ionization
states exhibit more of an asymmetric, slanted flat-topped shape. The bottom
panel shows a comparison between the high ionization states of the IMEs. The
[S IV] 10.51 $\mu$m appears more flat-topped owing to the neighboring [Co II]
10.52 $\mu$m line peaked toward the red. The different profile shapes between
ionization states imply ionization stratification. Figure 5: Left: Fits to Co
emission lines. The vertical dashed red line shows that the peak of the IGE
distributions is highly redshifted at $\sim 3500$ km s-1. Neighboring IME
lines are dashed blue lines while neighboring Fe lines are dashed gray lines.
Center top: Similar to the Co fits, but for Fe (neighboring Co lines are
dashed grey). Right: Fits to the IME emission lines. Neighboring IGE lines are
shown in dashed red. The vertical dashed blue line shows the peak of the IME
distributions, which is highly blueshifted at $\sim 2000$ km s-1. Center
bottom: Fits to the optical [O I] $\lambda\lambda 6300,~{}6364$ doublet and
the MIR [Ne II] 12.81 $\mu$m line. The [O I] is very central and narrow. The
[Ne II] peaks at the same location as the [O I] (vertical dashed green line);
however, it has an additional Gaussian-like red wing and a flat-topped-like
blue wing slanted in the direction opposite that of the IMEs.

###  3.2 Iron-Group Elements: Fe, Co, Ni

Starting with the isolated and well-resolved [Co III] 11.89 $\mu$m line, we
approximate the shape of the line profile by the sum of two Gaussians with
full width at half-maximum intensity (FWHM) $=4500$ km s-1and kinematic offset
$v_{\text{offset}}=4700$ km s-1, and FWHM $=9400$ km s-1and kinematic offset
$v_{\text{offset}}=1300$ km s-1. As shown in Figure 3, the [Fe II] 1.26 $\mu$m
line has a shape very similar (taking into account the neighboring [Fe II]
lines that create its red tail) to that of [Co III] 11.89 $\mu$m, so we fit
all the [Fe II] lines across 1.20–1.83 $\mu$m with a sum of two Gaussians and
obtain a similar, but slightly more peaked, profile for [Fe II] with FWHM
$=6700$ km s-1and kinematic offset $v_{\text{offset}}=3000$ km s-1, and FWHM
$=4700$ km s-1and kinematic offset $v_{\text{offset}}=-2500$ km s-1(see Figure
5 for a visualization of a double Gaussian profile).

A possible explanation for why Fe and Co do not exhibit exactly the same
profiles is that while Co mostly traces the original 56Ni profile, Fe is more
widely distributed because it is partially primordial. In the regions
dominated by the IMEs and O$+$Ne, Fe may be a significant coolant since
primordial Fe is $\sim 400$ times more abundant than Co. Additionally, or
alternatively, the Co may be less sharply peaked as this may be the region
corresponding to partial burning in quasistatistical equilibrium (DerKacy et
al., 2023).

Although the MIR isolated Ni lines are weak and low-resolution, [Ni II], [Ni
III], and [Ni IV] all exhibit a very similar line profile shape. Likewise, the
[Co II] lines have a similar shape as [Co III] and the [Fe III] lines have a
similar shape as [Fe II]. Thus, to reduce degeneracy and make fitting many
ions (and therefore parameters) more tractable, we use the same profile shape
parameters (excluding amplitude) for all ionization states of the same element
for Co, Fe, and Ni. We find that the Ni lines are slightly better fit by the
Fe parameters than by Co, so we use the same profile-shape parameters as [Fe
II] for all Ni lines. In principle, we would not necessarily expect different
ionization states of the same element to arise from the same emitting regions;
however, Figure 5 shows that our constraints on these profile fits can
reproduce the emission features very well from the optical through the MIR.
This suggests that there is not significant ionization stratification in the
IGEs (i.e., emission from different ionization states is coming from similar
ejecta regions). Furthermore, the IGE lines all peak with a redshift of $\sim
3500$ km s-1; the consistency of these profiles indicates that the IGEs are
present in roughly the same ejecta regions.

###  3.3 Intermediate-Mass Elements: S, Ar, Ca

The emission from IMEs in SN 2022pul is very strong and displays more line-
profile variation than the IGEs. Despite the variation, the IME emission-line
profiles share similar shapes and all exhibit a blueshifted peak at
$\sim-2000$ km s-1, the opposite direction as the IGEs. They also exhibit
ionization stratification, shown in Figure 4. The higher ionization states of
[S IV], [Ar III], and [Ca IV] all have a generally asymmetric boxy shape,
whereas the lower ionization [Ar II] and [Ca II] appear to have an additional
blueshifted, Gaussian-like component. We choose to fit the Ca II and Ar II
features with the sum of a Gaussian and a slanted flat-top profile. Physically
this could mean that the central, denser regions of the ejecta have lower
overall ionization, requiring this additional Gaussian component (see Figure 5
for a visualization of an asymmetric shell and additional Gaussian component
profile).

The right panel of Figure 5 shows our best-fit line profiles for the IMEs. The
[Ar III] 8.99 $\mu$m line is isolated and exhibits a clear slanted flat-topped
shape with Gaussian wings. This flat-topped Gaussian profile can arise from a
thick shell of [Ar III] emission, where one side of the shell is thicker than
the other, causing the slant. Viewing angle of this asymmetric shell also
contributes to the observed slant (DerKacy et al., 2023). The parameters for
our [Ar III] 8.99 $\mu$m fit are FWHM $=10,200$ km s-1 for the Gaussian wings,
overall kinematic offset $v_{\text{offset}}=800$ km s-1, inner shell velocity
$v_{\text{min}}=5500$ km s-1, and offset of the shell thickness
$v_{\text{offset, shell}}=\,-900$ km s-1.

As described above, we fix the fit parameters (excluding amplitude) for [Co
II] 10.52 $\mu$m and [Ni III] 11.00 $\mu$m, and then we fit the remaining
contribution to the feature at 10.5 $\mu$m as the [S IV] 10.51 $\mu$m line.
This yields a slanted boxy profile with Gaussian wings for [S IV], with a
similar width and slope as [Ar III], but a narrower inner shell radius. Its
parameters are FWHM $=\,9600$ km s-1, kinematic offset $v_{\text{offset}}=\,0$
km s-1, inner shell velocity $v_{\text{min}}=\,2900$ km s-1, and offset of the
shell thickness $v_{\text{offset, shell}}=\,-700$ km s-1.

When we assume the peak of the [Co II] 10.52 $\mu$m line is redshifted (like
all the relatively isolated IGE ions consistently display), we recover a
blueshifted peak for [S IV] that agrees with the other IMEs. Because the IGEs
are redshifted and the IMEs are blueshifted, we are able to distinguish the
contributions to the 10.5 $\mu$m feature from both [Co II] 10.52 $\mu$m and [S
IV] 10.51 $\mu$m. We do not attribute the full flat-topped, boxy shape of the
feature at 10.5 $\mu$m to [S IV] 10.51 $\mu$m alone because the neighboring
[Co III] 11.89 $\mu$m line is quite strong, and SN 2022pul has a low mean
ionization state, so we do not expect the contribution from [Co II] 10.52
$\mu$m to be negligible.

The [Ca IV] 3.21 $\mu$m line can also be fit reasonably well by an asymmetric
boxy profile with slightly larger width than [Ar III] and shallower slope.
Simultaneously fitting the nearby [Fe III] around 2.9 $\mu$m, we constrain the
contribution to the [Ca IV] profile from [Fe III] 3.23 $\mu$m to be very
minimal. Thus, most of the flux in this feature is from [Ca IV], though [Ni I]
3.12 $\mu$m may contribute to the narrow peak on the blue side of the [Ca IV]
emission.

[Ar II] 6.98 $\mu$m and [Ca II] $\lambda\lambda 7291,~{}7324$ show similar
line profiles (see Figure 3), and the [Ca II] line is helpful for interpreting
the general shape of the low-resolution [Ar II] line. Figure 4 shows that the
[Ar III] 8.99 $\mu$m line shape fits nicely with the [Ar II] 6.98 $\mu$m line
if a blueshifted component is added. Thus, we fit [Ar II] 6.98 $\mu$m by
summing a Gaussian with a forced contribution from the shape of the [Ar III].
Since the [Ca II] shares a similar shape to [Ar II], with a Gaussian-like peak
and a somewhat boxy red wing, we also fit it with the sum of a Gaussian and a
forced contribution from the shape of the [Ar III]. For [Ar II] the additional
Gaussian component has FWHM $=\,9100$ km s-1and a kinematic offset
$v_{\text{offset}}=\,-2200$ km s-1 and for [Ca II] the additional Gaussian
component has FWHM $=\,4600$ km s-1and a kinematic offset
$v_{\text{offset}}=\,-1600$ km s-1. Because the optical 0.73 $\mu$m region is
complex with contribution from many lines, including [Fe II], [Ni II], and [Ar
III], we fit all of these additional lines simultaneously with the [Ca II]
doublet and find their contribution to be relatively small. While we cannot
definitively rule out other line profile shapes for [Ca II] and [Ar II], it is
clear that they exhibit an additional blueshifted peak of emission not seen in
the higher IME ionization states.

The shapes of these IME line profiles and their ionization stratification
point to the presence of broad, asymmetric shells of higher ionization IME
material, and a somewhat narrower, blueshifted region of lower ionization IME
material.

Figure 6: Top left, top right, bottom left: Cartoon projected 2D geometric
interpretations of the IGE, IME, and O$+$Ne line profiles, assuming axial
symmetry. Bottom right: Bolometric luminosity per projected velocity bin (500
km s-1) along the line of sight for the IGEs (Fe, Co, Ni), the IMEs (S, Ar,
Ca), and O$+$Ne. There are clear differences in the overall line profiles from
each element group, indicating distinct geometric distributions of material.
This bolometric luminosity plot, derived from the spectral line profiles,
gives the true relative contributions of the different element groups without
suggesting microscopic mixing.

###  3.4 O and Ne

In the center-bottom panel of Figure 5, we display our best-fit line profiles
for [O I] and [Ne II]. The [O I] $\lambda\lambda 6300,~{}6364$ feature is
remarkably narrow (FWHM $\approx 2000$ km s-1), allowing us to resolve
emission from both components of the doublet. We do not detect emission from
either [O II] $\lambda$3727 or [O III] $\lambda\lambda$4959, 5007; however,
this region of the spectrum is dominated by low-ionization-state IGEs,
complicating the interpretation of the blended features. The [O I] emission is
best modeled by a single skewed Gaussian velocity component. Similar skewed
(or “sawtoothed”) emission has been observed in the [Fe II] and [Ca II]
features of other 03fg-like SNe Ia (Taubenberger, 2017; Siebert et al., 2023a,
b). This asymmetry could potentially be explained by either an intrinsically
asymmetric ejecta distribution or dust extinction from within the SN ejecta.

If a large amount of dust were present in the ejecta, assuming it did not
render the [O I] emission invisible, one would expect the redshifted emission
to be more heavily extinguished than the blueshifted emission. This would lead
to an emission profile whose total offset is preferentially blueshifted,
appears steeper on the blueshifted side, and is more shallow-sloped on the
redshifted side (Jerkstrand, 2017). There is some evidence that the [Ca II]
features of 03fg-like SNe are preferentially blueshifted (Siebert et al.,
2023b); however, one would expect more dust attenuation at shorter
wavelengths. Furthermore, the [O I] in SN 2022pul is redshifted by $\sim 300$
km s-1, so the systemic velocity of the ejecta (and dust) would need to be
redshifted if dust extinction were the cause of the asymmetric line profile.
This could be the case; however, we do not observe evidence for such a
systemic shift in the other nebular features, and therefore favor an
asymmetric ejecta distribution.

The [Ne II] 12.81 $\mu$m emission has a unique profile compared to the other
elements. There is a narrow peak close to $\sim 0$ km s-1, indicating that it
is present in the central regions of the ejecta along with [O I]. However,
most of the [Ne II] emission is likely not coincident with the [O I] owing to
its much larger ionization potential. The overall broad [Ne II] line suggests
that it forms at large velocities similar to the IMEs and IGEs. We model the
broad component of the [Ne II] line profile with a broad Gaussian on the red
side, and an asymmetric slanted flat-top on the blue side. The one-sided
asymmetric slant is tilted in the opposite direction as the tilt of the IMEs,
and could arise from a half-shell or some type of missing bubble of [Ne II] on
the blueshifted side (see Figure 5 for visualization of this profile). The
overall shape of this profile does not appear as blueshifted as the other
IMEs, and the differing locations of missing emission indicate that the
distributions of Ne and the IMEs are distinctly different.

The highly asymmetric [Ne II] line shares some of the same projected
velocities, but it cannot coexist with the IMEs and IGEs, indicating large-
scale composition asymmetry. Additionally, the small-scale structure seen
throughout the [Ne II] profile may be caused by the presence of a large number
of randomly distributed blobs (e.g., see Jerkstrand, 2017, Fig. 3);
alternatively, it may be due to higher noise levels at the longest MIRI/LRS
wavelengths. In Section 5, we discuss the detection of [Ne II] 12.81 $\mu$m in
relation to predictions from WD–WD violent merger models.

###  3.5 Geometric Interpretation: Distinct Ejecta Distributions

The difference in line profiles suggests differences in the geometric
distributions of material in the ejecta. In order to more easily compare
between the bulk emission coming from each group of elements (IGEs: Fe, Co,
Ni; IMEs: S, Ar, Ca; O, Ne), rather than from particular elements and ions, we
sum all of the line profiles across our full wavelength range within these
three groups in velocity space and calculate the bolometric luminosity along
the line of sight in bins of 500 km s-1 ($d\nu$ $L_{\text{bol}}$). Each
velocity bin represents a slab of emission on the plane of the sky, giving a
2D projected spatial distribution of the elements. Shown in the bottom right
panel of Figure 6, the summed profiles of the IGEs, IMEs, and O+Ne are close
to those of the individual profiles, but they erase the small variations seen
from profile to profile.

Using the geometric interpretations of the line-profile components that we
adopt in our fits, we construct cartoons for the projected 2D distribution of
our three element groups to aid visualization of the distinct regions (see
Figure 6). Assuming spherical symmetry, Gaussian emission distributions
produce Gaussian line profiles, and slanted flat-topped (or boxy) profiles
result from asymmetric shells of emission. We assume spherical symmetry for
the individual components of each profile (resulting in overall axial symmetry
for the 2D projected distributions) and stick to combinations of these two
basic line-profile shapes, so our cartoons are oversimplifications. The real
ejecta are probably asymmetric and complex, potentially including bubbles,
mushroom shapes, and clumps. Pakmor et al. (2012) show the distribution of
elements in their WD violent merger hydrodynamical simulation in their Fig. 2.
Interestingly, our cartoons bear some similarities to these physically
realistic distributions.

We stress that these cartoons are not to be taken literally and that they are
2D projections. Overlap in line-of-sight velocity ($v_{\text{LOS}}$) between
the element groups does not necessarily imply that they physically overlap in
3D space. Specifically, all of our distributions show emission at
$v_{\text{LOS}}=0$ km s-1; however, this is a projection effect because, from
theoretical models, there should be no IGEs or IMEs in the most central
regions where the [O I] emission is located.

Figure 7: Comparison of the dust and MIR-continuum-subtracted spectrum of SN
2022pul at 338 d post explosion to the DDT (blue), DOUBLEDET (green), GCD
(red), and MERGER (yellow) models from Blondin et al. (2023) (all recomputed
out to 338 d) in the 0.7–2 $\mu$m and 10–14 $\mu$m regions. The MERGER model
matches the observed spectrum of SN 2022pul best, notably producing stronger
[Ca II] near 0.73 $\mu$m, capturing the NIR [Fe II] lines extremely well, and
predicting the strong observed [Ne II] 12.81 $\mu$m line that is not seen in
the other models. However, there are also deficiencies to the MERGER model,
such as the strength of the Ni lines being overpredicted (e.g., [Ni III] 11.00
$\mu$m). This figure highlights differences between the model predictions; a
version covering the full wavelength range is given in the Appendix as Figure
10.

###  3.6 Nickel-56 Mass

At the late phase of SN 2022pul (338 d post explosion), the 56Ni (half-life of
$\sim 6$ d) has long since decayed away and the energy deposition that
continues to power the spectral emission lines comes from the radioactive
decay of 56Co (half-life of $\sim 77$ d). Locally-deposited gamma-rays and
positrons from the decay can be absorbed by the IGEs while the IMEs and
unburned elements can absorb nonlocally deposited gamma-rays and secondary,
low-energy photons. Thus, to estimate the initial mass of 56Ni synthesized in
the explosion, we account for the total flux by integrating the spectrum
($F=\int F_{\nu}\,d\nu$) in the range 0.4–14 $\mu$m. We then estimate the
luminosity from the emission lines by correcting the flux to a distance of 16
Mpc (see Paper I, Siebert et al. 2023b, for a discussion of distance to SN
2022pul). In the low-density nebular phase, however, most of the energy
escapes. To account for this, we divide by the fraction of the total decay
energy that is absorbed by the ejecta from the MERGER model (Blondin et al.,
2023, and described further in Section 4), which has a value
$E_{\text{absorbed}}/E_{\text{total}}=0.055$.

We estimate a bolometric luminosity (i.e., energy-leakage-corrected 56Co-decay
input power) at 338 d post explosion of $3.6\text{\times}{10}^{41}$ erg s-1,
which corresponds to a mass of synthesized 56Ni of 0.5 $M_{\odot}$(see
equations from Nadyozhin, 1994; Branch & Wheeler, 2017, Section 5.4). This
estimate suffers from uncertainties in distance ($\pm 2$ Mpc), reddening,
absolute flux calibration, gamma-ray escape fraction, and assumption of
spherical symmetry, but it is in rough agreement with the 0.6  $M_{\odot}$of
56Ni produced in the violent merger simulation by Pakmor et al. (2012) and the
calculated 56Ni masses for SN 2011fe, the canonical normal SN Ia (Bora et al.,
2022, and references therein).

##  4 WD–WD Violent Merger Model

Figure 8: Spectra of our reference (blue) and modified (red) merger models at
338 days post-explosion compared to SN 2022pul over selected wavelength ranges
(from top to bottom): optical (0.35–1 $\mu$m), NIR (1–5 $\mu$m), and MIR (5–14
$\mu$m). The SN 2022pul spectrum has been corrected for redshift and Milky Way
(MW) extinction. The synthetic fluxes correspond to the assumed distance to SN
2022pul of 16 Mpc; they have not been rescaled or normalized in any way. The
$\mathcal{F}_{X-Y\text{$\mu$m}}$ label gives the fraction of the total optical
to MIR flux for SN 2022pul emitted within the wavelength range of each plot.
We include line identifications based on their Sobolev equivalent width (see
Table 2 of Blondin et al. 2023). Transitions connected to the ground state are
marked with a “†” symbol. For consecutive lines within a multiplet (connected
by a horizontal line), we give the $gf$-weighted mean wavelength of the
transitions. Lines that only appear in one model are labeled with the
corresponding color. Figure 9: Top left: Density profiles at 338 d for the
MERGER model and our modified merger model with clumping and 0.1 $M_{\odot}$of
mass added to the central ($\lesssim 2000$ km s-1) region. Bottom left:
Temperature profiles. Top right: Mean Fe ionization profiles. The mean
ionization is defined as $\Sigma_{i}in^{i+}/\Sigma_{i}n^{i+}$, where $n^{i+}$
is the number density of ionization stage $i$ for Fe, such that a mean
ionization of $\sim 2$ indicates that Fe2+ is the dominant stage. Middle
right: Mean Ni ionization profiles. Bottom right: Mean O ionization profiles.

Blondin et al. (2023) presented nebular-phase spectra from four different SN
Ia explosion models at 270 d post-explosion (a $M_{\text{Ch}}$ delayed-
detonation model [DDT; Seitenzahl et al. 2013], a $M_{\text{Ch}}$
pulsationally-assisted gravitationally-confined detonation model [GCD; Lach et
al. 2022], a sub-$M_{\text{Ch}}$ double-detonation [DBLEDET; Gronow et al.
2021], and a violent merger of two sub-$M_{\text{Ch}}$ WDs (0.9$\,+\,$1.1
$M_{\odot}$) [MERGER; Pakmor et al. 2012]), which they compared to a spectrum
of SN 2021aefx covering the wavelength range 0.3–14 $\mu$m from Kwok et al.
(2023). These model spectra are generated in one-dimension, so the line
profile shapes do not reflect ejecta asymmetries that might be seen in two-
dimensional models. Only the MERGER model displayed a strong, centrally-peaked
line due to [Ne II] 12.81 $\mu$m, resulting from the presence of Ne in the
innermost ejecta layers of this model. None of the other models considered by
Blondin et al. (2023) presented such a strong [Ne II] feature, and only the
GCD model displayed a weak flat-topped line, whose intensity at line center
was a factor of $\sim 3$ smaller than observed in SN 2021aefx. Only by
artificially setting a minimum Ne mass fraction of $10^{-2}$ in the inner
ejecta of the GCD model were they able to match the peak intensity and profile
shape of the potential weak [Ne II] line in SN 2021aefx. Future analysis of
subsequent observations of SN 2021aefx should clarify whether [Ne II] 12.81
$\mu$m is actually present in SN 2021aefx (Ashall et al., in prep.).

The violent merger model naturally predicts the presence of Ne in the inner
ejecta layers, which is associated with the delayed and incomplete burning of
the lower-mass secondary WD. Moreover, the higher density in layers $\lesssim
10,000$ km s-1 enhances the recombination rate there, which in turn results in
a lower mean ionization state compared to the other explosion models. For
comparing to the nebular spectrum of SN 2022pul, we recomputed the four models
studied by Blondin et al. (2023) at a later time of 338 d post-explosion,
using the same numerical setup. Highlights of the results comparing these
models to the observed spectrum are displayed in Figure 7, and the full
optical through MIR comparison can be found in the Appendix. Of the four
models, only the MERGER model is consistent with the strong and centrally-
peaked [Ne II] 12.81 $\mu$m line.

The lower ionization results in a better match to the numerous [Fe II]
features in the 1–2 $\mu$m range, yet the optical [Fe III] complex around 0.5
$\mu$m is still too strong compared to SN 2022pul, as is the [Ni III] 7.35
$\mu$m line. This suggests the ionization in SN 2022pul is even lower than
predicted in the MERGER model. We consider clumping to be a viable mechanism
to further reduce the ionization state in this model. Using the same approach
as Blondin et al. (2023), we recomputed a clumped version of the MERGER model
with a volume-filling factor $f=0.5$ (i.e., resulting in a factor of two
increase in density of the clumps). As seen in Figure 8, this clumped model is
able to better match the flux level in the [Fe III]-dominated complex, while
maintaining a satisfactory match out to $\sim 3$ $\mu$m. The lower ionization
also results in a stronger [Ar II] 6.98 $\mu$m line and a weaker [Ni III] 7.35
$\mu$m line. However, it also results in an even weaker [Ca IV] 3.12 $\mu$m
line compared to SN 2022pul. The same is true of [Ar III] 8.99 $\mu$m.
Moreover, the [Ni II] 6.63 $\mu$m and [Ni III] 7.35 and 11.00 $\mu$m lines
remain too strong. This suggests that the clumping is not uniform, affecting
some regions more strongly than others.

The mismatch in lines from specific elements over several ionization stages
suggests that abundance also plays a role. Here, a lower abundance of stable
Ni and a higher abundance of Ar/Ca could help resolve some of these
discrepancies. In particular, a lower abundance of stable Ni ($\sim 0.03$
$M_{\odot}$ in this MERGER model) seems warranted to reproduce the lines of Ni
II and Ni III in the observed spectrum. Pakmor et al. (2013) show in their
helium-ignited violent merger of a system of 0.9$\,+\,$0.76  $M_{\odot}$ WDs
that the central density and the Ar/Ca abundance is higher in the innermost
layers, while the stable Ni abundance is lower and more centrally concentrated
than the merger from Pakmor et al. (2012), which the MERGER model is based on.
Additionally, the amount of stable Ni in the merger model from Pakmor et al.
(2012) scales with the metallicity of the progenitor system ($Z=Z_{\odot}$ in
the model), so lower stable Ni abundance may imply the progenitor system had
subsolar metallicity. There may be a variety of inner-ejecta properties that
could be produced through violent merger scenarios that would better match the
observations of SN 2022pul.

Despite the lower ionization of the clumped MERGER model, the [O I]
$\lambda\lambda 6300,~{}6364$ emission remains very weak compared with SN
2022pul. This is not related to an abundance issue, since oxygen completely
dominates the ejecta layers below $\sim 2000$ km s-1 in the MERGER model (mass
fraction $\gtrsim 0.5$). Radial compression does not change the optical depth,
so the mass of inner material must be increased to boost the amount of gammy-
ray energy intercepted by the inner, O-rich ejecta layers. We were able to
match the intensity and width of the narrow [O I] feature seen in SN 2022pul
by artificially increasing the density below $\sim 2000$ km s-1. In practice
we followed the procedure from Dessart et al. (2020) and added a Gaussian
component to the original density profile centered on the inner boundary of
our spatial grid ($v_{\mathrm{min}}\approx 580$ km s-1) with a width
$\sigma\approx 600$ km s-1, resulting in an additional 0.1  $M_{\odot}$ in
these layers ($\sim 50$% of which is O). The original mass of O (0.53
$M_{\odot}$) was all located in the inner ejecta and the total ejecta mass in
the model is increased through this modification. The fractional O abundance
in the innermost region is not particularly important, as long as [O I]
$\lambda\lambda 6300,~{}6364$ remains the dominant coolant.

As seen from Figure 8, the narrow [O I] feature is well reproduced, while the
rest of the spectrum remains largely unchanged compared to the original model.
The physical origin of such an increase in density in the inner ejecta is
likely the result of a higher-mass secondary WD (e.g., Pakmor et al. 2013).
Other contributing effects could be how completely the secondary WD burned, a
different C/O fraction of the initial secondary WD, and compression of the
innermost layers through the 56Ni bubble effect (e.g., Wang, 2005; Dessart et
al., 2021). The changes to the ejecta properties (density, temperature, mean
ionization) in our modified merger model where clumping and 0.1 $M_{\odot}$
central mass have been introduced are shown in Figure 9, and a detailed list
of important emission lines in our modified merger model is given in Table 1.

##  5 Discussion & Conclusions

As discussed in Paper I (Siebert et al., 2023b), the detection of [O I]
emission in WD SNe is exceedingly rare. SN 2022pul is now the third 03fg-like
SN Ia to exhibit [O I], along with SN 2012dn (Taubenberger et al., 2019) and
SN 2021zny (Dimitriadis et al., 2023). The velocity distribution of [O I]
varies greatly among these SNe, while its offset is consistent with
originating from the central regions of the ejecta in each case. Additionally,
SN 2022pul exhibits [Ne II] peaked in the central regions and extending
farther out as well. A viable explosion model should be able to explain the
presence and diversity of oxygen and neon. Both turbulent pure deflagrations
(Kozma et al., 2005; Fink et al., 2014) and WD mergers (Pakmor et al., 2012;
Kromer et al., 2013; Pakmor et al., 2013) naturally predict neon and unburned
oxygen present at low velocities. A WD–WD merger provides an explanation for
the large-scale asymmetries in the nebular emission features described in
Section 3, though a turbulent deflagration could also produce asymmetries.
However, a pure deflagration would produce a low-luminosity explosion with
different early-time spectroscopic properties, narrower nebular line profiles
owing to the lower explosion energy, and weaker emission from IMEs. To check
this, we ran the N100def model from Fink et al. (2014) with CMFGEN to 338 d
post explosion and indeed found that the MERGER model is a significantly
better match to SN 2022pul for the above reasons. We therefore favor a WD
merger over a pure deflagration for SN 2022pul.

The large parameter space in WD mergers could help explain the diversity of
03fg-like SNe. The variation among objects may be explained by different WD
mass ratios, by the specific explosion mechanism of one or both WDs, and by
observations along different lines of sight. Recent studies of 03fg-like SNe
Ia have favored WD merger models occurring within a dense C/O-rich CSM
(Dimitriadis et al., 2022; Srivastav et al., 2023; Dimitriadis et al., 2023;
Siebert et al., 2023a) due to rapidly evolving light-curve bumps, low-
ionization nebular spectra sometimes with [O I], asymmetric and blueshifted
nebular emission lines, and a larger fraction of flux in the NIR at late times
than normal SNe Ia. Our observations of SN 2022pul enhance the evidence for
this model for 03fg-like SNe Ia through clear dust detection, unusual light-
curve rise shape, asymmetric line profiles, strong IMEs, and central O and Ne.

In the case of SN 2022pul, we also observe distinct ejecta components with
different composition offset in their bulk velocities. The secondary WD may
have been partially, but not completely, disrupted when the primary WD
detonated. When burning front passes over the lower density secondary WD
material, it too can burn, synthesizing O, Ne, and the IMEs (whereas the
higher density primary WD would produce the IGEs) (Pakmor, 2017). The degree
of ejecta asymmetry will be smaller if the secondary is more compact when the
primary is burned because the ashes of the primary will be less hindered by a
less disrupted secondary (Pakmor, 2017).

In SN 2022pul, we observe central O and Ne, very strong emission from IMEs,
and a high degree of asymmetry in the nebular line features. Furthermore, our
modified merger model demonstrates that additional, concentrated central mass
(0.1 M⊙ at $\lesssim 2000$ km s-1) is needed to reproduce the [O I]
$\lambda\lambda 6300,~{}6364$ emission. This very centrally peaked and
asymmetric [O I] emission is reminiscent of expectations of stripped material
from a companion (Marietta et al., 2000; Botyánszki et al., 2018), and may
indicate that the secondary WD left a significant amount of mass in the
innermost region.

The central, broad Ne and strong IMEs suggest that the secondary WD material
was at least partially burned (the O may be a combination of burned C and
unburned O, while unburned Ne is not expected for the lower-mass secondary). A
higher mass, and therefore less disrupted secondary WD, could produce the
strong IMEs and the centrally concentrated material. However, the high degree
of asymmetry would suggest a more disrupted secondary, which perhaps could
have left behind centrally concentrated material if it was only partially
burned. This may be due in part to a viewing-angle effect where the ejecta
appear more symmetric in the vertical direction, and more asymmetric in the
plane of rotation (Bulla et al., 2016; Pakmor, 2017). Previous measurements of
03fg-like SNe Ia have not shown strong continuum polarization (Tanaka et al.,
2010; Cikota et al., 2019). More polarization data and construction of nebular
spectra along different viewing angles from violent merger hydrodynamical
simulations would help resolve this and improve our understanding of the 3D
ejecta distributions giving rise to the asymmetric line profiles.

Another key part of understanding this progenitor scenario is its connection
to the inferred presence of a dense C/O-rich CSM from the shape of the light-
curve rise (Siebert et al., 2023b). This requires considering scenarios where
the explosion occurs either pre-merger during a dynamical interaction phase
(commonly known as a “perimerger” or violent merger), or scenarios where
explosion occurs post-merger after the complete disruption of the secondary WD
(Pakmor et al., 2010; van Kerkwijk et al., 2010; Pakmor et al., 2012; Moll &
Woosley, 2013; Guillochon et al., 2010; Raskin & Kasen, 2013).

In both cases, simulations of the ejecta are highly asymmetric and the
synthetic light curves cover a large range in luminosity depending on viewing
angle (with the highest luminosity observed along the equator). This light-
curve line-of-sight dependence might also vary over time because at early
phases the interaction dominates over decay power. One beneficial consequence
of the post-merger scenario is that the complete disruption of the secondary
WD leads to the formation of a massive C/O-rich accretion disk. The early-time
light curve of SN 2022pul is consistent with ejecta interaction with this kind
of CSM (Noebauer et al., 2016; Siebert et al., 2023b). More work is needed to
understand if a violent merger is capable of producing enough CSM at the
required distance to explain this kind of interaction signature.

The large amount of CSM (0.6 $M_{\odot}$) adopted in the Noebauer et al.
(2016) interaction model is difficult to reconcile with the central oxygen and
neon seen in SN 2022pul. If this much material from the secondary WD was
pulled off into an exterior spherical shell, it seems unlikely that the
secondary would have enough leftover mass to create the high-density central O
giving rise to the [O I] doublet and the strong emission from the IMEs. If the
CSM is responsible for damping the kinetic energy of the explosion, leading to
the boosted luminosity and low-velocity C II at early times, the CSM must
cover a substantial fraction of the exploding system (Maeda et al., 2023).
Howell et al. (2006) suggest that the extra binding energy associated with the
additional mass may be enough to slow down the ejecta sufficiently, and the
amount of CSM stripped from the secondary WD in the Raskin & Kasen (2013)
models is much smaller, less than 0.005 $M_{\odot}$. If the dust emission
arises from CSM, assuming a dust-to-gas ratio of 0.01, the dust mass for SN
2022pul given by Johansson et al. (2023) implies a CSM mass of 0.05–0.11
$M_{\odot}$, depending on composition. Further exploration is needed to
address whether these lower CSM masses (or perhaps a different composition)
could produce the interaction signature we observe.

Hsiao et al. (2020) present evidence supporting the “core-degenerate” model
for LSQ14fmg, where a near-$M_{\text{Ch}}$ WD interacts with the degenerate
core of an AGB star within a common envelope stripped of its H and He through
a superwind phase prior to explosion. We disfavor this scenario to explain SN
2022pul because the explosion is expected to be a spherical explosion of a
single object after secular accretion. It therefore does not explain the
observed central location of O and Ne or the large-scale ejecta asymmetries.
It is tempting to speculate whether a modified version of this model, in which
the AGB degenerate core material also undergoes thermonuclear fusion, could be
developed for SN 2022pul.

Perhaps a triple system could be invoked (Kushnir et al., 2013; Grishin &
Perets, 2022), where one WD or star is disrupted to become the surrounding CSM
and precipitates the violent merger of the remaining two WDs. Rajamuthukumar
et al. (2023) found that triple systems contribute significantly to the
overall SN Ia rate. Again, more detailed modeling over a wide range of
possible parameters and scenarios involving WD–WD violent mergers is needed to
better understand the system giving rise to this unique SN.

We summarize our main observations and discussion points below.

  * •

The presence of O and Ne in the central ejecta are hallmarks of the WD violent
merger SN Ia model.

  * •

Asymmetric nebular emission-line profiles show distinct distributions of IGEs,
IMEs, and O and Ne in the ejecta that can be naturally explained by the
violent merger model.

  * •

A pure deflagration model is disfavored as it would produce lower
luminosities, narrower nebular line widths, and weaker emission from IMEs than
observed in SN 2022pul.

  * •

Especially strong emission from IMEs could come from burning of a disrupted,
lower-density secondary WD.

  * •

Our spectrum of SN 2022pul is quite well matched by the MERGER model nebular
spectrum computed by Blondin et al. (2023) from the violent WD–WD merger model
of Pakmor et al. (2012). The agreement between the observed and model spectra
is further improved by the addition of clumping and 0.1 $M_{\odot}$ to the
central mass.

  * •

It is still somewhat unclear how the violent merger model can produce both the
C/O-rich CSM described by Paper I (Siebert et al., 2023b) and Paper III
(Johansson et al., 2023), and the central O and Ne seen in the spectra.

  * •

Further work and more detailed models are needed to constrain the masses of
the WDs involved in the merger.

Our observations of SN 2022pul are a prime example of the importance of
gathering spectra of WD SNe across the optical, NIR, and MIR wavelengths. The
MIR observation of [Ne II] is complementary to the observation of [O I] in the
optical; the optical [Ca II], NIR [Ca IV], and MIR [Ar II], [Ar III], and [S
IV] mutually help interpret the IME distribution and ionization
stratification; and the NIR [Fe II] and MIR [Co III] lines establish the
distribution of the IGEs. Furthermore, the higher resolution ground-based
optical and NIR observations were essential to interpret the lower resolution
JWST observations, and isolated MIR spectral lines were key to unraveling the
complex line blending in the optical and NIR.

Future observations of 03fg-like SNe Ia would benefit from JWST spectroscopy,
especially in the higher resolution instrument modes which reach to even
longer wavelengths and reveal additional lines. Complementary
spectropolarimetry would give insight into line-profile asymmetries and
constrain aspect angles (DerKacy et al., 2023). If this class of objects
indeed arises from WD mergers, the observations should show a large degree of
diversity in their asymmetric line profiles, and JWST observations could
additionally establish a firmer connection to the dust surrounding or formed
in these objects.

This work is based on observations made with the NASA/ESA/CSA JWST as part of
program #02072. We thank Shelly Meyett for her consistently excellent work
scheduling the JWST observations, Sarah Kendrew for assistance with the MIRI
observations, and Glenn Wahlgren for help with the NIRSpec observations.

The data were obtained from the Mikulski Archive for Space Telescopes at the
Space Telescope Science Institute (STScI), which is operated by the
Association of Universities for Research in Astronomy (AURA), Inc., under
National Aeronautics and Space Administration (NASA) contract NAS 5-03127 for
JWST. Support for this program at Rutgers University was provided by NASA
through grant JWST-GO-02072.001.

The SALT observations of SN 2022pul were obtained with Rutgers University
program 2022-1-MLT-004 (PI S. W. Jha). We are grateful to SALT Astronomer
Rosalind Skelton for taking these data.

This work makes use of data from the Las Cumbres Observatory global network of
telescopes. The LCO group is supported by NSF grants AST-1911151 and
AST-1911225. This work also makes use of data gathered with the 6.5 meter
Magellan telescopes at Las Campanas Observatory, Chile.

M.R.S. is supported by an STScI Postdoctoral Fellowship. G.D. acknowledges
H2020 European Research Council grant #758638. L.A.K. acknowledges support by
NASA FINESST fellowship 80NSSC22K1599. C.L. is supported by an NSF Graduate
Research Fellowship under grant # DGE-2233066.

The UCSC team is supported in part by NASA grant NNG-17PX03C, National Science
Foundation (NSF) grant AST-1815935, the Gordon and Betty Moore Foundation, the
Heising-Simons Foundation, and a fellowship from the David and Lucile Packard
Foundation to R.J.F. The work of A.V.F.’s supernova group at UC Berkeley is
generously supported by the Christopher R. Redlich Fund, Alan Eustace (W.Z. is
a Eustace Specialist in Astronomy), Frank and Kathleen Wood (T.G.B. is a Wood
Specialist in Astronomy), Gary and Cynthia Bengier, Clark and Sharon Winslow,
Sanford Robertson (Y.Y. is a Bengier-Winslow-Robertson Postdoctoral Fellow),
and many other donors.

S.B. acknowledges support from the Alexander von Humboldt Foundation and from
the “Programme National de Physique Stellaire” (PNPS) of CNRS/INSU cofunded by
CEA and CNES. D.J.H. receives support through NASA astrophysical theory grant
80NSSC20K0524. J.V. and T.S. are supported by the NKFIH-OTKA grants K-142534
and FK-134432 of the Hungarian National Research, Development and Innovation
(NRDI) Office, respectively. B.B. and T.S. are supported by the ÚNKP-22-4 and
ÚNKP-22-5 New National Excellence Programs of the Ministry for Culture and
Innovation from the source of the NRDI Fund, Hungary. T.S. is also supported
by the János Bolyai Research Scholarship of the Hungarian Academy of Sciences.
The research of J.C.W. and J.V. is supported by NSF grant AST-1813825.

A.F. acknowledges support by the European Research Council (ERC) under the
European Union’s Horizon 2020 research and innovation program (ERC Advanced
Grant KILONOVA #885281). M.D., K.M., and J.H.T. acknowledge support from EU
H2020 ERC grant #758638. Research by Y.D. and S.V. is supported by NSF grant
AST-2008108

Time-domain research by D.J.S. and the University of Arizona team is supported
by NSF grants AST-1821987, 1813466, 1908972, and 2108032, and by the Heising-
Simons Foundation under grant #2020–1864.

K.M. acknowledges support from the Japan Society for the Promotion of Science
(JSPS) KAKENHI grant JP20H00174, and the JSPS Open Partnership Bilateral Joint
Research Project (JPJSBP120209937).

L.G. acknowledges financial support from the Spanish Ministerio de Ciencia e
Innovación (MCIN), the Agencia Estatal de Investigación (AEI)
10.13039/501100011033, and the European Social Fund (ESF) “Investing in your
future” under the 2019 Ramón y Cajal program RYC2019-027683-I and the
PID2020-115253GA-I00 HOSTFLOWS project, from Centro Superior de
Investigaciones Científicas (CSIC) under the PIE project 20215AT016, and the
program Unidad de Excelencia María de Maeztu CEX2020-001058-M.

J.P.H. acknowledges support from the George A. and Margaret M. Downsbrough
bequest. The Aarhus supernova group is funded in part by an Experiment grant
(#28021) from the Villum FONDEN, and by a project 1 grant (#8021-00170B) from
the Independent Research Fund Denmark (IRFD).

This publication was made possible through the support of an LSSTC Catalyst
Fellowship to K.A.B., funded through grant 62192 from the John Templeton
Foundation to the LSST Corporation. The opinions expressed in this publication
are those of the authors and do not necessarily reflect the views of LSSTC or
the John Templeton Foundation.

Some of the data presented herein were obtained at the W. M. Keck Observatory,
which is operated as a scientific partnership among the California Institute
of Technology, the University of California, and NASA. The Observatory was
made possible by the generous financial support of the W. M. Keck Foundation.
The authors wish to recognize and acknowledge the very significant cultural
role and reverence that the summit of Maunakea has always had within the
indigenous Hawaiian community. We are most fortunate to have the opportunity
to conduct observations from this mountain.

A major upgrade of the Kast spectrograph on the Shane 3 m telescope at Lick
Observatory was made possible through generous gifts from the Heising-Simons
Foundation as well as William and Marina Kast. Research at Lick Observatory is
partially supported by a generous gift from Google.

We thank the Subaru staff for the data taken by the Subaru Telescope
(S23A-023).

Table 1: Lines in the wavelength range 0.35-14 $\mu$m in our modified merger model (MERGER with clumping + 0.1  $M_{\odot}$ central mass) at 338 days post explosion whose absolute Sobolev EW exceeds 5% of the largest absolute EW in the wavelength range 0.35-1 $\mu$m. $\lambda_{\mathrm{air}}$ | Ion | $\lambda_{\mathrm{air}}$ | Ion | $\lambda_{\mathrm{air}}$ | Ion | $\lambda_{\mathrm{air}}$ | Ion | $\lambda_{\mathrm{air}}$ | Ion | $\lambda_{\mathrm{air}}$ | Ion | $\lambda_{\mathrm{air}}$ | Ion
---|---|---|---|---|---|---|---|---|---|---|---|---|---
($\mu$m) |  |  ($\mu$m) |  |  ($\mu$m) |  |  ($\mu$m) |  |  ($\mu$m) |  |  ($\mu$m) |  |  ($\mu$m) |
0.427 |  Fe i a | 0.501 |  [Fe iii] i | 0.717 |  [Fe ii] m | 0.923 |  [Fe ii] t | 1.321 |  [Fe ii] y | 1.954 |  [Fe ii] aa | 5.672 |  [Fe ii] ah
0.429 ${\dagger}$ |  [Fe ii] b | 0.501 |  Fe i h |  0.729 ${\dagger}$ |  [Ca ii] n | 0.927 |  [Fe ii] t | 1.328 |  [Fe ii] y | 1.967 |  [Fe ii] ab | 5.703 |  [Co ii] ai
0.431 |  Fe i a | 0.502 |  Fe ii g |  0.732 ${\dagger}$ |  [Ca ii] n |  0.934 ${\dagger}$ |  [Co ii] v | 1.372 |  [Fe ii] y | 2.007 |  [Fe ii] ad | 6.212 |  [Co ii] ai
0.433 |  Fe i a | 0.505 |  Fe i h |  0.738 ${\dagger}$ |  [Ni ii] o | 0.934 |  [Co ii] w | 1.497 |  [Co ii] z | 2.015 |  [Fe ii] ae |  6.634 ${\dagger}$ |  [Ni ii] a
0.436 |  [Fe ii] b | 0.508 |  Fe i h | 0.739 |  [Fe ii] m | 0.953 |  [S iii] u | 1.533 |  [Fe ii] aa | 2.046 |  [Fe ii] ad | 6.719 |  [Fe ii] ah
0.438 ${\dagger}$ |  Fe i] c | 0.516 |  [Fe ii] a | 0.741 |  [Ni ii] o | 0.964 |  [Co ii] v | 1.547 |  [Co ii] z | 2.133 |  [Fe ii] ad | 6.918 |  [Ni ii] a
0.438 |  Fe i] d | 0.517 |  Fe i j | 0.745 |  [Fe ii] m | 0.994 |  [Co ii] w | 1.599 |  [Fe ii] aa | 2.206 |  Na i af |  6.983 ${\dagger}$ |  [Ar ii] a
0.440 |  Fe i] d | 0.517 |  Fe ii g | 0.764 |  [Fe ii] p |  1.019 ${\dagger}$ |  [Co ii] v | 1.626 |  [Co ii] z | 2.208 |  Na i af |  7.347 ${\dagger}$ |  [Ni iii] aj
0.441 |  [Fe ii] b | 0.517 |  Fe i a | 0.769 |  [Fe ii] p | 1.025 |  [Co ii] v | 1.634 |  [Co ii] z | 2.224 |  [Fe ii] ae |  7.505 ${\dagger}$ |  [Ni i] ak
0.442 |  Fe i] d | 0.523 |  Fe i j | 0.789 |  [Ni iii] q | 1.028 |  [Co ii] v | 1.644 |  [Fe ii] aa | 2.244 |  [Fe ii] ad | 7.788 |  [Fe iii] a
0.442 ${\dagger}$ |  [Fe ii] a | 0.527 |  Fe i k | 0.833 |  Fe i r | 1.032 |  [S ii] a | 1.664 |  [Fe ii] aa | 2.308 |  [Ni ii] ac |  8.403 ${\dagger}$ |  [Ni iv] al
0.443 |  Fe i] c | 0.527 |  [Fe iii] i | 0.839 |  Fe i r | 1.097 |  [Co ii] v | 1.677 |  [Fe ii] aa | 2.369 |  [Ni ii] ac |  8.989 ${\dagger}$ |  [Ar iii] a
0.445 |  [Fe ii] b | 0.533 |  Fe i k | 0.850 |  Ca ii s | 1.128 |  [Co ii] v | 1.711 |  [Fe ii] aa | 2.911 |  [Ni ii] ac |  10.508 ${\dagger}$ |  [S iv] a
0.446 |  Fe i] c | 0.537 |  Fe i k | 0.850 |  [Ni iii] q | 1.161 |  Fe i x | 1.728 |  [Co ii] z | 3.119 |  [Ni i] a |  10.520 ${\dagger}$ |  [Co ii] v
0.448 |  Fe i] c | 0.540 |  Fe i k | 0.854 |  Ca ii s | 1.169 |  Fe i x | 1.736 |  [Co ii] z |  3.206 ${\dagger}$ |  [Ca iv] a | 10.679 |  [Ni ii] am
0.455 |  Fe ii e | 0.541 |  Fe i k | 0.862 |  [Fe ii] t | 1.188 |  Fe i x | 1.745 |  [Fe ii] aa | 3.393 |  [Ni iii] ag | 10.999 |  [Ni iii] aj
0.458 |  Fe ii e | 0.543 |  Fe i k | 0.866 |  Fe i r | 1.188 |  Fe i x | 1.748 |  [Fe ii] ab | 3.801 |  [Ni iii] ag | 11.164 |  [Co ii] v
0.466 ${\dagger}$ |  [Fe iii] f | 0.543 |  Fe i k | 0.866 |  Ca ii s | 1.197 |  Fe i x | 1.797 |  [Fe ii] aa | 4.075 |  [Fe ii] ah | 11.304 |  [Ni i] ak
0.470 |  [Fe iii] f | 0.545 |  Fe i k | 0.869 |  Fe i r | 1.249 |  [Fe ii] y | 1.800 |  [Fe ii] aa |  4.114 ${\dagger}$ |  [Fe ii] ah | 11.723 |  [Ni iv] al
0.481 |  [Fe ii] a | 0.546 |  Fe i k | 0.882 |  Fe i r |  1.257 ${\dagger}$ |  [Fe ii] y | 1.809 |  [Fe ii] aa | 4.434 |  [Fe ii] ah |  11.885 ${\dagger}$ |  [Co iii] a
0.489 |  [Fe ii] a |  0.589 ${\dagger}$ |  Na i l | 0.889 |  [Fe ii] t | 1.270 |  [Fe ii] y | 1.811 |  [Fe ii] ab | 4.606 |  [Fe ii] ah | 11.998 |  [Ni i] a
0.492 |  Fe ii g |  0.590 ${\dagger}$ |  Na i l | 0.903 |  [Fe ii] t | 1.279 |  [Fe ii] y | 1.895 |  [Fe ii] aa | 4.888 |  [Fe ii] ah |  12.252 ${\dagger}$ |  [Co i] a
0.494 |  Fe i h |  0.630 ${\dagger}$ |  [O i] a | 0.905 |  [Fe ii] t | 1.294 |  [Fe ii] y | 1.903 |  [Co ii] z | 5.061 |  [Fe ii] ah | 12.725 |  [Ni ii] am
0.499 |  Fe i h | 0.716 |  [Fe ii] m | 0.907 |  [S iii] u | 1.298 |  [Fe ii] y | 1.939 |  [Ni ii] ac |  5.339 ${\dagger}$ |  [Fe ii] ah |  12.811 ${\dagger}$ |  [Ne ii] a

Figure 10: Spectra of our reference model set at 338 days post explosion
compared to SN 2022pul over the wavelength ranges (from top to bottom) optical
(0.35–1 $\mu$m), NIR (1–5 $\mu$m), lower MIR (5–14 $\mu$m), and upper MIR
(14–28 $\mu$m). The SN 2022pul spectrum has been corrected for redshift and MW
extinction. The synthetic fluxes correspond to the assumed distance to SN
2022pul of 16 Mpc; they have not been rescaled or normalized in any way. The
$F_{X-Y\text{$\mu$m}}$ label gives the fraction of the total optical to MIR
flux (0.35–14 $\mu$m) for SN 2022pul emitted within the wavelength range of
each plot. We include line identifications based on their Sobolev equivalent
width (see Table 2 of Blondin et al. 2023). Transitions connected to the
ground state are marked with a “†” symbol. For consecutive lines within a
multiplet (connected by a horizontal line), we give the $gf$-weighted mean
wavelength of the transitions.

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