# JWST NIRSpec+MIRI Observations of the nearby Type IIP supernova 2022acko

M. Shahbandeh Space Telescope Science Institute, 3700 San Martin Drive,
Baltimore, MD 21218-2410, USA  C. Ashall Department of Physics, Virginia Tech,
Blacksburg, VA 24061, USA  P. Hoeflich Department of Physics, Florida State
University, 77 Chieftan Way, Tallahassee, FL 32306, USA  E. Baron Planetary
Science Institute, 1700 East Fort Lowell Road, Suite 106, Tucson, AZ
85719-2395, USA  Hamburger Sternwarte, Gojenbergsweg 112, D-21029 Hamburg,
Germany  Dept of Physics & Astronomy, University of Oklahoma, 440 W. Brooks,
Rm 100, Norman, OK USA  O. Fox Space Telescope Science Institute, 3700 San
Martin Drive, Baltimore, MD 21218-2410, USA  T. Mera Department of Physics,
Florida State University, 77 Chieftan Way, Tallahassee, FL 32306, USA  J.
DerKacy Department of Physics, Virginia Tech, Blacksburg, VA 24061, USA  M. D.
Stritzinger Department of Physics and Astronomy, Aarhus University, Ny
Munkegade 120, DK-8000 Aarhus C, Denmark  B. Shappee Institute for Astronomy,
University of Hawai’i at Manoa, 2680 Woodlawn Dr., Hawai’i, HI 96822, USA  D.
Law Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD
21218-2410, USA  J. Morrison Space Telescope Science Institute, 3700 San
Martin Drive, Baltimore, MD 21218-2410, USA  T. Pauly Space Telescope Science
Institute, 3700 San Martin Drive, Baltimore, MD 21218-2410, USA  J. Pierel
Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD
21218-2410, USA  K. Medler Astrophysics Research Institute, Liverpool John
Moores University, UK  J. Andrews Gemini Observatory/NSF’s NOIRLab, 670 North
A‘ohoku Place, Hilo, HI 96720-2700, USA  D. Baade European Organization for
Astronomical Research in the Southern Hemisphere (ESO), Karl-Schwarzschild-
Str. 2, 85748 Garching b. München, Germany  A. Bostroem LSSTC Catalyst Fellow
Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson,
AZ 85721-0065, USA  P. Brown George P. and Cynthia Woods Mitchell Institute
for Fundamental Physics and Astronomy, Texas A&M University, Department of
Physics and Astronomy, College Station, TX 77843, USA  C. Burns Observatories
of the Carnegie Institution for Science, 813 Santa Barbara Street, Pasadena,
CA 91101, USA  A. Burrow Homer L. Dodge Department of Physics and Astronomy,
University of Oklahoma, 440 W. Brooks, Rm 100, Norman, OK 73019-2061, USA  A.
Cikota Gemini Observatory/NSF’s NOIRLab, Casilla 603, La Serena, Chile  D.
Cross  Institute of Space Sciences (ICE, CSIC), Campus UAB, Carrer de Can
Magrans, s/n, E-08193 Barcelona, Spain  S. Davis Private Astronomer  T. de
Jaeger LPNHE, (CNRS/IN2P3, Sorbonne Université, Université Paris Cité),
Laboratoire de Physique Nucléaire et de Hautes Énergies, 75005, Paris, France
A. Do Institute for Astronomy, University of Hawai’i at Manoa, 2680 Woodlawn
Dr., Hawai’i, HI 96822, USA  Y. Dong Department of Physics, University of
California, 1 Shields Avenue, Davis, CA 95616-5270, USA  I. Dominguez
Universidad de Granada, 18071, Granada, Spain  L. 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  D. Janzen Department of Physics and Engineering Physics,
University of Saskatchewan, 116 Science Place, Saskatoon, SK S7N 5E2, Canada
J. Jencson Department of Physics and Astronomy, The Johns Hopkins University,
3400 North Charles Street, Baltimore, MD 21218, USA  E. Hoang Department of
Physics and Astronomy, University of California, Davis, 1 Shields Avenue,
Davis, CA 95616-5270, USA  E. Hsiao Department of Physics, Florida State
University, 77 Chieftan Way, Tallahassee, FL 32306, USA  E. Karamehmetoglu
Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120,
DK-8000 Aarhus C, Denmark  B. Khaghani Department of Physics, Virginia Tech,
Blacksburg, VA 24061, USA  K. Krisciunas George P. and Cynthia Woods Mitchell
Institute for Fundamental Physics and Astronomy, Texas A&M University,
Department of Physics and Astronomy, College Station, TX 77843, USA  S. Kumar
Department of Physics, Florida State University, 77 Chieftan Way, Tallahassee,
FL 32306, USA  J. Lu Department of Physics, Florida State University, 77
Chieftan Way, Tallahassee, FL 32306, USA  M. Lundquist W. M. Keck Observatory,
65-1120 Māmalahoa Highway, Kamuela, HI 96743-8431, USA  J. Maund Department of
Physics and Astronomy, University of Sheffield, Hicks Building, Hounsfield
Road, Sheffield S3 7RH, U.K.  P. Mazzali Astrophysics Research Institute,
Liverpool John Moores University, UK  Max-Planck Institute for Astrophysics,
Garching, Germany  N. Morrell Las Campanas Observatory, Carnegie
Observatories, Casilla 601, La Serena, Chile  F. Patat European Organization
for Astronomical Research in the Southern Hemisphere (ESO), Karl-
Schwarzschild-Str. 2, 85748 Garching b. München, Germany  J. Pearson Steward
Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ
85721-0065, USA  C. Pfeffer Department of Physics, Virginia Tech, Blacksburg,
VA 24061, USA  M. Phillips Las Campanas Observatory, Carnegie Observatories,
Casilla 601, La Serena, Chile  A. Rest  Space Telescope Science Institute,
3700 San Martin Drive, Baltimore, MD 21218-2410, USA  N. Retamal Department of
Physics, University of California, 1 Shields Avenue, Davis, CA 95616-5270, USA
S. Stangl Homer L. Dodge Department of Physics and Astronomy, University of
Oklahoma, 440 W. Brooks, Rm 100, Norman, OK 73019-2061, USA  M. Shrestha
Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson,
AZ 85721-0065, USA  C. Stevens Department of Physics, Virginia Tech,
Blacksburg, VA 24061, USA  N. Suntzeff George P. and Cynthia Woods Mitchell
Institute for Fundamental Physics and Astronomy, Texas A&M University,
Department of Physics and Astronomy, College Station, TX 77843, USA  C.
Telesco Department of Astronomy, University of Florida, Gainesville, FL 32611
USA  M. Tucker CCAPP Fellow  Center for Cosmology and AstroParticle Physics,
The Ohio State University, 191 W. Woodruff Ave., Columbus, OH 43210, USA  L.
Wang Department of Physics and Astronomy, Texas A&M University, College
Station, TX 77843, USA  Y. Yang Department of Astronomy, University of
California, Berkeley, CA 94720-3411, USA  Bengier-Winslow-Robertson
Postdoctoral Fellow  Y.-Z. Cai  Yunnan Observatories, Chinese Academy of
Sciences, Kunming 650216, PR China  Key Laboratory for the Structure and
Evolution of Celestial Objects, Chinese Academy of Sciences, Kunming 650216,
PR China  International Centre of Supernovae, Yunnan Key Laboratory, Kunming
650216, P.R. China  Y. Camacho-Neves Department of Physics and Astronomy,
Rutgers, the State University of New Jersey,
136 Frelinghuysen Road, Piscataway, NJ 08854-8019, USA  N. Elias-Rosa INAF -
Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5, 35122, Padova,
Italy  R. Foley Department of Astronomy and Astrophysics, University of
California, Santa Cruz, CA 95064-1077, USA  S. Jha Department of Physics and
Astronomy, Rutgers, the State University of New Jersey,
136 Frelinghuysen Road, Piscataway, NJ 08854-8019, USA  L. Kwok Department of
Physics and Astronomy, Rutgers, the State University of New Jersey,
136 Frelinghuysen Road, Piscataway, NJ 08854-8019, USA  C. Larison Department
of Physics and Astronomy, Rutgers, the State University of New Jersey,
136 Frelinghuysen Road, Piscataway, NJ 08854-8019, USA  N. LeBaron  Department
of Astronomy, University of California, Berkeley, CA 94720-3411, USA  S. Moran
Department of Physics and Astronomy, University of Turku, Vesilinnantie 5,
FI-20500, Finland  J. Rho  SETI Institute, 339 N. Bernardo Ave., Ste. 200,
Mountain View, CA 94043, USA  Department of Physics and Astronomy, Seoul
National University, Gwanak-ro 1, Gwanak-gu, Seoul, 08826, South Korea  I.
Salmaso Dipartimento di Fisica e Astronomia “G. Galilei”, Università degli
Studi di Padova, Vicolo dell’Osservatorio 3, 35122, Padova, Italy  INAF -
Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5, 35122, Padova,
Italy  J. Schmidt  Private Astronomer  S. Tinyanont National Astronomical
Research Institute of Thailand, 260 Moo 4, Donkaew, Maerim, Chiang Mai 50180,
Thailand

(Received XXX; Accepted xxx)

###### Abstract

We present JWST spectral and photometric observations of the Type IIP
supernova (SN) 2022acko at $\sim$50 days past explosions. These data are the
first JWST spectral observations of a core-collapse SN. We identify $\sim 30$
different H i features, other features associated with products produced from
the CNO cycle, and s-process elements such as Sc ii, Ba ii. By combining the
JWST spectra with ground-based optical and NIR spectra, we construct a full
Spectral Energy Distribution from 0.4 to 25 $\micron$ and find that the JWST
spectra are fully consistent with the simultaneous JWST photometry. The data
lack signatures of $CO$ formation and we estimate a limit on the $CO$ mass of
$<10^{-8}M_{\odot}$. We demonstrate how the $CO$ fundamental band limits can
be used to probe underlying physics during the stellar evolution, explosion,
and the environment. The observations indicate little mixing between the H
envelope and C/O core in the ejecta and show no evidence for dust. The data
presented here set a critical baseline for future JWST observations, where
possible molecular and dust formation may be seen.

supernovae: general - supernovae: individual (SN 2022acko), JWST

††journal: ApJL††facilities: JWST (MRS/MIRI,NIRSpec/MIRI), MAST
(JWST)††software: jwst (Bushouse et al., 2023, ver. 1.9.4,), jdaviz
(Developers et al., 2023, ver. 3.2.1,), HYDRA (Höflich, 2003, 2009; Hoeflich
et al., 2017), OpenDx (an open-sourced visualization package developed by
IBM), Astropy (Astropy Collaboration et al., 2013, 2018, 2022), NumPy (Harris
et al., 2020), SciPy (Virtanen et al., 2020), Matplotlib (Hunter, 2007),
pypeit (Prochaska et al., 2020).

##  1 Introduction

The origin of the substantial quantities of dust inferred from submillimeter
and millimeter observations of high-redshift ($z$) quasars is an unsolved
astrophysical question (e.g., Bertoldi et al., 2003; Priddey et al., 2003; Li
et al., 2020). In the local Universe, most dust is believed to originate
during the evolutionary phase of low- and intermediate-mass stars when they
reach the asymptotic giant branch on the Hertzsprung-Russell diagram (e.g.,
González Delgado et al., 2003; Ferrarotti & Gail, 2006), However, at higher
redshifts (i.e., $z>6.3$), alternative processes occurring on shorter time
scales are more likely at play.

For more than half a century, core-collapse supernovae (CCSNe) have been
proposed as a potential source of dust (Cernuschi et al., 1967; Hoyle &
Wickramasinghe, 1970). In particular, hydrogen-rich CCSNe known as Type II
plateau supernovae (hereafter, SNe IIP), are linked to the demise of massive
stars born with initial masses ranging between 8 to 25 $M_{\odot}$, and with
lifetimes of less than 50 Myrs. SNe IIP have played a significant role in the
evolution of the Universe, contributing significantly to the production of
heavy elements, and the genesis of dust, both locally and at cosmologically
significant redshifts (e.g., Maiolino et al., 2004; Schneider et al., 2004;
Dwek et al., 2007; Nozawa et al., 2008; Gall et al., 2011). Models of the
expanding SN ejecta have demonstrated their ability to condense sufficient
quantities (0.1–1 $M_{\sun}$) of dust (e.g., Todini & Ferrara, 2001; Nozawa et
al., 2003; Cherchneff & Dwek, 2009; Sarangi & Cherchneff, 2015; Sluder et al.,
2018; Sarangi et al., 2018).

Only a handful of SNe IIP have been documented to form dust, and until
recently, these cases produced 2–3 orders of magnitude less dust than
predicted by modeling. The prerequisite for the genesis of dust is the
creation of molecules, which serve both as a coolant and as a nucleation site.
In SNe IIP, the order of the formation of silicon-monoxide ($SiO$) versus
carbon-monoxide ($CO$) can be dependent upon the initial zero-age main
sequence (ZAMS) mass of the progenitor stars since it leads to differing
amounts of C, O, and Si in the stars at the time of their demise (Woosley et
al., 2002).

The majority of previous studies on molecular formation in CCSNe have
concentrated on searching for the ($\approx 2.3~{}\micron$) $CO$ first
overtone feature. Evidence of molecular formation was first found in data of
SN 1987A (e.g., Catchpole et al., 1988; Meikle et al., 1989; Wooden et al.,
1993), and to date has also been identified in around a dozen SNe IIP (e.g.,
Aitken et al., 1988; Kotak et al., 2005a; Meikle et al., 2007a; Kotak et al.,
2009; Fox et al., 2010; Tinyanont et al., 2019; Szalai et al., 2019; Davis et
al., 2019).

The ability to observe the $CO$ and $SiO$ fundamental bands began to improve
with the advent of the Spitzer Space Telescope (SST), as the longer spectral
wavelength range was particularily well-suited for studying molecular
formation in CCSNe. SST spectra of SN 2005af showed indications of $SiO$
formation (Kotak et al., 2006), while observations of SN 2004dj contained the
red part of the $CO$ fundamental band (Kotak et al., 2005b). Given the nature
of the data, these initial assessments of molecules in a handful of SNe left
open the question of whether we see pre-existing or freshly synthesized dust
(Gerardy et al., 2000).

The James Webb Space Telescope (JWST) has revolutionized the communities
ability to accurately study molecule and dust formation in various types of
SNe. Recently, Shahbandeh et al. (2023) used JWST Mid-Infrared Instrument
(MIRI; Bouchet et al. 2015; Ressler et al. 2015; Rieke et al. 2015, 2022)
observations of the Type IIP SN 2004et to infer a dust mass lower limit of
$0.014$ $M_{\odot}$. This exceeds estimates estimated from Mid-infrared (MIR)
observations of all other observed SNe IIP (see Shahbandeh et al., 2023) , and
highlights the potential of SNe IIP being a major contributor to the dust
content of the Universe.

While deep imaging is possible, JWST spectroscopy offers additional insights
into the explosion physics of CCSNe. Near-infrared (NIR) and MIR spectroscopic
time series covering the duration of a CCSN’s lifetime can provide information
on the formation of molecules like $CO$ and $SiO$. To fully understand how SNe
IIP produces molecules and dust, we must follow them beginning soon after
exploding and over the duration of many decades. Where early data allows us to
set a baseline for any pre-existing molecular and dust formation.

We present Near-Infrared Spectrograph (NIRSpec; Jakobsen et al. 2022) and MIRI
observations of the Type IIP SN 2022acko, as part of JWST GO-2122 (PI: C.
Ashall), which aims to follow the evolution of a single SN IIP from the early
H recombination driven plateau phase thru late times when $CO$ and $SiO$
molecules form. Here, we present the first NIRSpec+MIRI spectral observation
from this program, the first CCSN spectrum obtained with JWST. We also present
near-simultaneous optical and NIR ground-based spectral observations of SN
2022acko, as well as serendipitous multi-band imaging obtained by the Physics
at High Angular Resolution of Nearby GalaxieS (PHANGS) JWST Cycle 1 Treasury
Program (Lee et al., 2023).

Our goals are to i) measure the level of mixing within the ejecta by searching
for $CO$ formation, trace the elements of the CNO-cycle, and identify possible
s-process elements, ii) set limits for pre-existing molecules and dust in the
interstellar medium (ISM) and circumstellar medium (CSM), and iii) demonstrate
the importance of fundamental bands of molecules, obtain values and tight
upper limits for the nearby environment, for mixing processes during the
stellar evolution, and for the physics of the explosion. This is a key to
being able to separate pre-existing molecules and dust from newly formed
products, and to evaluate the SN IIP as dust-producers in the early Universe.

SN 2022acko was discovered on 2022 December 06.2 UT (i.e., MJD-59919.2) by the
Distance Less Than 40 Mpc Survey (DLT40; Tartaglia et al. 2018), and a
spectrum obtained 12 hours with the Lijiang 2.4-m telescope showed it to be a
young SN IIP (Li et al., 2022). With J2000 coordinates R.A.
$=03^{h}19^{m}38^{s}.99$ and Decl. $=-19^{\circ}23^{\prime}42\farcs 68$, the
location of SN 2022acko is $0.0181^{\circ}$ North-West from the center of NCG
1300 (de Vaucouleurs et al., 1991). According to the NASA extragalactic
database (NED) the heliocentric redshift ($z$) to the host-galaxy of SN
2022acko is $z=0.00526$. In the following, we adopt the tip of red giant
branch (TRGB) distance measurement from the PHANGS team of $18.99\pm 2.85$ Mpc
(Anand et al., 2021).

Its proximity as well as the location of SN 2022acko outside the core of its
host made it an ideal target for our JWST-GO-2122 ToO program (PI: C. Ashall,
DOI:10.17909/ea3g-5z06). Ultraviolet and optical observations of SN 2022acko
indicate it is a low-luminosity SN IIP (see Bostroem et al., 2023). For
consistency, throughout this work, we adopt the same basic parameters for SN
2022acko as inferred by Bostroem et al. (see Table 1), including their
preferred reddening values and estimated explosion epoch which they
demonstrated likely occurred within 24 hours prior to discovery.

Table 1: SN 2022acko details.  Parameter | Value | Reference
---|---|---
RA | 03:19:38.99 | Lundquist et al. (2022)
Dec |  $-$19:23:42.68 | Lundquist et al. (2022)
Redshift | 0.00526 | Bostroem et al. (2023)
Host galaxy | NGC 1300 | Lundquist et al. (2022)
$E(B-V)_{MW}$ | 0.026 $\pm$0.001 mag | Bostroem et al. (2023)
$E(B-V)_{host}$ |  $0.04\pm 0.01$ mag | Bostroem et al. (2023)
Distance [Mpc] | $19.0\pm 2.9$ | Anand et al. (2021)
Explosion epoch ($t_{exp}$; MJD)aafootnotemark:  | $59918.17\pm 0.4$ | Bostroem et al. (2023)

The outline of this paper is as follows. Our observations of SN 2022acko are
presented in Section 2, which is then followed by the identification of key
spectral features and data comparisons in Section 3. Next in Section 4, H line
velocities are measured, after which the full optical-to-MIR spectral energy
distribution (SED) constructed in Section 5. Limits on pre-existing molecular
and dust formation are determined in Section 6, and then in Section 7, SN
2022acko is compared with SN 1987A in order to understand the amount of mixing
in the ejecta. Finally, a summary of the results and our conclusions are given
in Section 8.

##  2 Observations

###  2.1 NIRCam and MIRI Imaging

PHANGS obtained imaging of NGC 1300 with NIRCam and MIRI on 59969.74 MJD
(51.57 days post-explosion), which is within $\sim 1$ day of our programme’s
NIRSpec observations of SN 2022acko (see below). This data was acquired in
four NIR filters (i.e., F200W, F300M, and F335M, F360M), as well as in four
MIR filters (i.e., F770W, F1000W, F1130W, and F2100W) providing a measure of
SN 2022acko’s Spectral Energy Distribution (SED) between $\sim$1.7–21.0
$\mu$m. Calibrated level 2 data products were downloaded from the Mikulski
Archive for Space Telescopes (MAST), from which point-spread function (PSF)
photometry was computed.

A color-stacked image constructed from MIRI images is presented in Fig. 1. The
position of SN 2022acko at the top of the image is indicated with a red circle
while a $10\times 10$ arcsec zoom of this region is provided in the bottom-
right inset. SN 2022acko is located in a region of star formation. Figure 2
shows the SN in a 30 by 30 pixel stamp in the $F300M$, $F355M$, $F360M$,
$F770W$, $F1000W$ and $F1130W$ filters. The SN is clearly visible in each of
these bands.

PSF photometry was computed using a coustom-made phython notebook making use
of the WebbPSF photometry package (Perrin et al.,
2014).111https://github.com/orifox/psf_phot/blob/main/space_phot/MIRI/miri_1028.ipynb
Various PSF sizes were used to ensure that the full variation within the data
was encapsulated. The final magnitude was taken as the average flux across all
fits, which have a residual less than 15%. The fluxes of all four dithers of
each filter were then averaged, and the error in the flux was obtained using
the standard deviation of the fluxes. A log of the photometric observations
and the flux at these epochs, can be found in Table 2.

Figure 1: A color image of NGC 1300 constructed from three JWST MIRI images. The color code of the images are is followed, Red: F2100W, Orange: F1130W, Cyan: F770W. SN 2022acko is located in the top of the image in the red circle. A 10” by 10” inset around the location of the SN is seen in the bottom right. The SN is located in an area of star formation. The full resolution of this image can be found here. Figure 2: NIRCAM F300M, F355M, F360M 30 pixel $\times$ 30 pixel stamps of SN 2022acko, and MIRI F770W, F1000W and F1130W 30 pixel $\times$ 30 pixel stamps of SN 2022acko. We do not show the F2100W filter as the SN is not visible.  Table 2:  Log of JWST photometric observations. The average time of photometric observation is 51.6 d after the inferred time of the explosion.  Filters | Obs. date | Exp. time | Mag |
---|---|---|---|---
| [MJD] | [s] | [mag] |
F200Waafootnotemark:  | 59969.82 | 215 | 17.00 $\pm$ 0.09ccfootnotemark:  |
F300Maafootnotemark:  | 59969.79 | 387 | 17.53 $\pm$ 0.03 |
F335Maafootnotemark:  | 59969.80 | 387 | 17.60 $\pm$ 0.05 |
F360Maafootnotemark:  | 59969.77 | 215ddfootnotemark:  | 17.62 $\pm$ 0.01 |
F770Wbbfootnotemark:  | 59969.65 | 88.8 | 18.63 $\pm$ 0.01 |
F1000Wbbfootnotemark:  | 59969.66 | 122 | 19.19 $\pm$ 0.01 |
F1130Wbbfootnotemark:  | 59969.66 | 311 | 19.35 $\pm$ 0.01 |
F2100Wbbfootnotemark:  | 59969.68 | 322 |  $\leq$21.43 $\pm$ $\cdots$ |

###  2.2 NIRSpec

We obtained a NIRSpec spectrum (R$\sim$ 1000) of SN 2022acko from
$\sim$1.7-5.2 $\micron$. The spectrum was acquired using the S400A1 slit, the
G235M and G395M gratings, with the F170P and F290LP filters, respectively. The
instrument setup had one exposure per grating and a total exposure time of
467.65 s. The details of the observational setup can be seen in Table 3. The
data were taken on MJD 59968.4 at 50.2 d past explosion. The observation
consisted of a three-point dither pattern. The background-subtracted level-two
spectra were combined into a single calibrated spectrum for each filter. The
spectra were processed with the JWST Calibration Pipeline version 1.8.2, using
the Calibration Reference Data System version 11.16.14 (Bushouse et al.,
2023).

Table 3: SN 2022acko spectral log.

Parameter | Value | Value | Value
---|---|---|---
MIRI/MRS Spectra
Grating | Short | Medium | Long
Groups per Integration | 36 | 36 | 36
Integrations per Exp. | 1 | 1 | 1
Exposures per Dither | 1 | 1 | 1
Total Dithers | 4 | 4 | 4
Exp Time [s] | 3440.148 | 3440.148 | 3440.148
Resolutionaafootnotemark:  |  $\sim$2,700
Epochaafootnotemark:  [days] | 55.57
$T_{\text{obs}}$ [MJD] | 59973.74
NIRSpec Acquisition Image
Filter | CLEAR | $\cdots$ | $\cdots$
Exp Time [s] | 0.08 | $\cdots$ | $\cdots$
Readout Pattern | NRSRAPID | $\cdots$ | $\cdots$
NIRSpec Spectra
Slit | S400A1 | S400A1 | $\cdots$
Filter | F170P | F290LP | $\cdots$
Groups per Integration | 5 | 7 | $\cdots$
Integrations per Exp. | 2 | 2 | $\cdots$
Exposures per Dither | 1 | 2 | $\cdots$
Total Dithers | 3 | 3 | $\cdots$
Exp Time [s] | 196.43 | 271.22 | $\cdots$
Epochaafootnotemark:  [days] | 50.18 | 50.19 | $\cdots$
$T_{\text{obs}}$ [MJD] | 59968.35 | 59968.36 | $\cdots$
NIR Spectra/Keck-NIRES
Exp time [s] | 1200 | $\cdots$ | $\cdots$
$T_{\text{obs}}$ [MJD] | 59951.95 | $\cdots$ | $\cdots$
Epochaafootnotemark:  [days] | 33.78 | $\cdots$ | $\cdots$
Total Dithers | 4 | $\cdots$ | $\cdots$
Silt width | 0.5” | $\cdots$ | $\cdots$
Optical Spectra/NOT-ALFOSC
Exp time [s] | 1800 | $\cdots$ | $\cdots$
$T_{\text{obs}}$ [MJD] | 59978.86 | $\cdots$ | $\cdots$
Epochaafootnotemark:  [days] | 60.69 | $\cdots$ | $\cdots$
Grism | Grism_#4 | $\cdots$ | $\cdots$
Silt width | 1.0” | $\cdots$ | $\cdots$

Note. — aafootnotemark: NIRCAM data, bbfootnotemark: MIRI data. All
photometric observations were performed with a 4-point dither. ccfootnotemark:
The F200W images were saturated in all exposures, therefore the photometry is
highly uncertain. All exposure times are reported from four dithers unless
otherwise stated.

Note. — aafootnotemark:  Rest frame days relative to time of explosion, MJD
$=59918.17$ (Bostroem et al., 2023).

Figure 3: Top panel: MIRI/MRS cube of SN 2022acko before background
subtraction divided seperate by channel. Each channel is a collapsed sum of
all its slices. Bottom panel: MIRI/MRS cubes of SN 2022acko after background
subtraction.

Figure 4: First: NOT (+ ALFOSC) optical spectrum of SN 2022acko, with labeled
line IDS. Second: The Keck-II (+ NIRES) NIR spectrum with labeled line IDs.

Figure 5:  First: The JWST-NIRSpec S400A1/F170P spectrum with labeled line
IDs. Second: The JWST-NIRSpec S400A1/F290P spectrum with labeled line IDs.

Figure 6: The reduced stitch 4 channel MIRI-MRS spectrum of SN 2022acko with
line identifications. Line identifications in Channel 4 are tentative due to
the large background on the observations.

###  2.3 MIRI Medium Resolution Spectroscopy

Medium resolution spectroscopy (MRS) data were obtained with the MIRI
instrument on JWST. Observations were taken in the Short, Medium, and Long
gratings to ensure continuous wavelength coverage (5$-$25 $\micron$) and
stitched together. We note that the data longward of 15 $\micron$ is dominated
by noise and instrumental thermal background; this is in part because of the
scientific goals of these observations were to cover the $SiO$ fundamental
overtone at $\sim$8.1 $\mu$m, the exposure time was not optimized for the
longest wavelengths, and because of the known issues with the degrading count
rate throughput of Ch3 and
Ch4222https://www.stsci.edu/contents/news/jwst/2023/miri-mrs-reduced-count-
rate-update. The MRS data were reduced using a custom-built
pipeline333https://github.com/shahbandeh/MIRI_MRS designed to extract point
source observations with highly varying and/or complex backgrounds from MRS
data cubes. The pipeline creates a master background based on 48 different
positions in the data cube. This master background is then subtracted from the
whole data cube before the aperture photometry is performed along the data
cube, using the Extract1dStep in stage 3 of the JWST reduction pipeline. The
reduction shown here utilized version 1.12.1 of the JWST Calibration pipeline
(Bushouse et al., 2023) and Calibration Reference Data System files version
11.17.1. Figure 3 shows the stacked MRS data cubes in each channel before and
after background subtraction, and the spectrum is shown in full in Figure 6.
The raw data associated with this reduction can be found at 10.17909/ea3g-5z06
(catalog DOI:10.17909/ea3g-5z06).

###  2.4 Ground-based spectroscopy

To complete the SED to shorter wavelengths, we also present ground-based
spectra of SN 2022acko at similar epochs. A summary of these optical and NIR
observations is provided in Table 1.

An optical spectrum was obtained with the NUTS2 collaboration using the 2.54-m
Nordic Optical Telescope (NOT) equipped with the Alhambra Faint Object
Spectrograph and Camera (ALFOSC). The spectrum was taken on MJD 59978.86 which
corresponds to +60.7 days post-explosion. The data were reduced within the
PyRAF environment using the NUTS2 pipeline written by E. Cappellero.

A NIR spectrum was taken on MJD 59951.95 at +33.78 days post-explosion with
the Near-Infrared Echellette Spectrometer (NIRES) mounted to the 10-m Keck II
telescope. SN 2022acko was observed at two positions along the slit (AB pairs)
to perform background subtraction. An A0V star was observed immediately
adjacent to the science observation, and used for both telluric and flux
calibration. The NIRES data were reduced using the PYTHON-based pypeit
package. Combining these four spectra produces the first-ever nearly-
continuous spectral coverage between $\sim$ 0.4–25 $\micron$, which
demonstrates the power of combining ground-based and JWST data.

Two additional spectra were taken using the ALFOSC on the NOT via the NUTS
collaboration, and an additional ground-based NIR spectrum was taken with EMIR
mounted on the Gran Telescopio Canarias (GTC). These data are included in the
online supplementary material.

##  3 Line Identifications and Data Comparsion

The following section discusses the spectral lines contributing to the
spectral formation. The spectra of SN 2022acko are taken towards the end of
the plateau phase. The hydrogen (H) spectral series, including the Balmer,
Paschen, Brackett, Pfund, and Humphreys dominate the spectra of SN 2022acko.
We also see some features of heavier elements (e.g., Mazzali et al., 1992).
Table 4 lists all of the spectral lines that have been identified in the full
spectrum, and below, we discuss the line identifications within various
regions.

###  3.1 Optical Spectrum

The top panel of Fig. 4 shows the optical spectrum of SN 2022acko. For
assistance with line IDs we use the models from Baron et al. (2003). In the
optical SN 2022acko is dominated by the Balmer series with some Paschen
series: Hα (0.6563 µm), Hβ (0.4861 µm), Hγ (0.4340 µm), Hδ (0.4102 µm), Paη
(0.901 µm), Paζ (0.923 µm), and Paϵ (0.955 µm). Other ions which contribute to
the spectral line formation include: the Ca ii H&K lines (0.3933 µm, 0.3968
µm) and Ca ii NIR triplet (0.8498 µm, 0.8542 µm, 0.8662 µm), Mg ii 0.9218 µm,
Na i 0.5896 µm, 0.5890 µm, Sc ii 5527 µm, 5698 µm, 6280 µm, Ba ii 6142 µm, and
many lines of Fe ii including 0.4359 µm, 0.4642 µm, 0.4862 µm, 0.4924 µm,
0.5018 µm, and 0.5169 µm. As well as C I lines particularly the line at 0.94
µm, just redward of Paζ.

###  3.2 NIR ground-based spectrum

We use the lines IDs from Davis et al. (2019) and Shahbandeh et al. (2022) to
assist us with identifying the main lines which contribute to the ground-based
NIR spectrum. With the exception of Paα, which is located in the NIRSpec data,
the H features in the ground-based NIR spectrum are dominated by the Paschen
series. These include Paβ (0.1282 µm), Paγ (0.1094 µm), and Paδ (0.1005 µm).
Other lines identified in this spectrum are Sr ii 1.0040 µm, 1.0330 µm, O i
1.1290 µm, Mg ii1.0092, 1.0927 µm, Ca ii 1.1839 µm, 1.1950 µm, and C i 1.0694
µm, 1.1333 µm, 1.175 µm.

###  3.3 NIR JWST/NIRSpec spectrum

Figure 5 shows the NIRSpec spectrum obtained at $\sim$+50 d. The spectrum is
dominated by P Cygni profiles of H lines mainly coming from the Brackett and
Pfund series. The main lines in the spectra include Brϵ (1.817 µm), H i 1.875
µm, Brδ (1.944 µm), Brγ (2.166 µm), Brβ (2.626 µm), H i 2.873 µm, 3.297 µm,
3.741 µm, Brα (4.052 µm), H i 4.377 µm, Pfβ (4.653 µm), H i 5.129 µm.

###  3.4 MIR JWST/MIRI spectrum

Figure 6 shows the lines identified in the MIRI spectrum. The MIRI spectrum is
comprised of observations across four channels, which range from 4.90-7.65 µm
for Channel 1, 7.51-11.70 µm for Channel 2, 11.55-17.98 µm for Channel 3, and
17.70-27.90 µm for Channel 4. Visual inspection of the non-background
subtracted data cubes shows there is clear SN flux in the first two channels;
the singal-to-noise (S/N) of the SN varies between 20 and 2 from 4.90-11.70
µm, after background subtraction, this S/N increases to 35 to 6 (see Figure
3). This gives us confidence that our line identifications shortwards of 11.70
µm are accurate. Before background subtraction, in Channel 3, the S/N varies
from 1-2 depending on the exact wavelength range. After background
subtraction, the S/N is $\sim$3\. Due to the high background and low
throughput of MIRI-MRS in channel 4, visual inspection of the pre-background
subtracted data shows no sign of the SN. However, after background
subtraction, (i.e. when the data is fully reduced using our pipeline) the
supernova is visible in Channel 4 at a S/N of 1-3. For example, we clearly see
the SN at wavelength 22.167 $\micron$, 25.1854 $\micron$, 26.9623 $\micron$.
Therefore, we still choose to identify possible spectral lines in this region
but as we could not manually inspect every wavelength slice for possible SN
signal, we want to remind the reader that the identifications in this Channel
are tentative.

Generally, the MIR spectrum is dominated by H lines (transitions from levels
6, 7, 8, and 9), however, there are several features that are from other ions.
Some of the most notable spectral features include C ii 5.539, 5.836, 6.775,
6.900 $\micron$, N ii 8.439, N iii 9.175, 9.191 $\micron$, O i 7.720, 10.384
$\micron$ and O ii 9.332, 9.439 $\micron$. The observation of s-process
elements and lines of C, N, and O in the hydrogen elements validate our
understanding of CNO burning during stellar evolution.

Table 4: SN 2022acko spectral lines.  Ion | Wavelength | Ion | Wavelength | Ion | Wavelength
---|---|---|---|---|---
| [$\mu$m] |  | [$\mu$m] |  | [$\mu$m]
Optical spectrum
H i | 0.4861 | Sc iii | 0.6238 | Sc iii | 0.7449
H i | 0.3889 | H i | 0.6563 | H i | 0.9229
Ca ii | 0.3969 | Fe ii | 0.4924 | Ca ii | 0.8498
H i | 0.3970 | Fe ii | 0.5169 | Ca ii | 0.8498
Ca ii | 0.3934 | Sc ii | 0.5527 | Ca ii | 0.8542
H i | 0.4102 | Sc ii | 0.5698 | Ca ii | 0.8662
H i | 0.4341 | O i | 0.7774 | Mg ii | 0.9218
Fe ii | 0.4358 | Na i | 0.5894 | C i | 0.9408
Fe ii | 0.4642 | Ba ii | 0.6142 | H i | 0.9014
Ground-based NIR spectrum
H i | 1.005 | Mg ii | 1.0092 | O i | 1.1290
H i | 1.0938 | Mg ii | 1.0927 | Sr ii | 1.0330
H i | 1.2818 | Na i | 1.0749 | Sr ii | 1.0920
Ca ii | 1.1839 | C i | 1.1333 |  |
Ca ii | 1.1950 | C i | 1.1756 |  |
NIRSpec spectrum
H i | 1.817 | H i | 2.626 | H i | 4.052
H i | 1.875 | H i | 2.873 | H i | 4.377
H i | 1.944 | H i | 3.297 | H i | 4.654
H i | 2.166 | H i | 3.741 | H i | 5.129
MIRI/MRS spectrum
H i | 5.129 | O i | 10.384 | Fe i | 16.596
C ii | 5.539 | O i | 10.399 | H i | 16.881
H i | 5.711 | H i | 11.309 | H i | 18.615
C ii | 5.836 | O i | 11.853 | N i | 18.811
H i | 6.291 | O i | 12.058 | H i | 18.978
H i | 6.772 | H i | 12.372 | H i | 19.062
C ii | 6.775 | H i | 12.390 | N i | 19.616
C ii | 6.900 | O i | 12.995 | N i | 19.641
H i | 7.460 | C i | 13.297 | N i | 20.353
O i | 7.720 | C i | 13.370 | N i | 21.022
C i | 8.374 | C i | 13.872 | C i | 22.278
N ii | 8.439 | H i | 13.942 | H i | 23.868
H i | 8.760 | H i | 14.183 | N i | 25.058
N iii | 9.175 | H i | 14.962 | H i | 26.168
N iii | 9.191 | Fe i | 15.561 | H i | 26.682
O ii | 9.332 | C i | 15.803 | C i | 27.425
H i | 9.392 | Fe i | 16.107 | H i | 27.803
O ii | 9.439 | H i | 16.209 |  |

###  3.5 Spectral Comparison

In this section, we examine the differences between SN 2022acko and other CCSN
which have MIR spectra. Figure 7 shows a comparison of the Spitzer Space
Telescope data of SN 2005af obtained at day +67 (Kotak et al., 2006) with that
of SN 1987A at day +56 (Aitken et al., 1988) and the MIRI/MRS data of SN
2022acko obtained at day +56 d. The three SNe are generally similar but have a
few notable differences. At this epoch, SN 2022acko is still on the plateau
phase and is dominated by H features, whereas SN 2005af was faster evolving
and passed the plateau phase by day +67 (Kotak et al., 2005b). In fact, the
forbidden lines identified by Kotak et al. (2006) are absent in both SN 1987A
and SN 2022acko. Although the SN 2022acko MIRI spectrum is $\sim 10$ days
earlier than SN 2005af, it appears to be evolving more slowly, thus we would
expect the ZAMS mass of SN 2022acko to be larger than that of SN 2005af. SN
2004dj was also observed by the SST at epochs of 89–129 days (Kotak et al.,
2005b). In this spectra [Ni ii] lines were seen and the presence of $CO$ was
inferred as early as 106 days. Future observations of SN 2022acko will allow
us to compare the epoch of $CO$ formation, and evolution of the onset of
forbidden lines, which will enable us to determine if SN 2022acko had a larger
ejecta mass.

Figure 7: Comparison between MIR spectra of SN 1987A, 2005af and 2022acko. the
forbidden lines seen in the day 67 spectrum of SN 2005af are absent in earlier
spectra of SN 1987A and 2022acko. This is because at these phases SN 2005af
was at the end of the plateau. the data from SN 1987A are digitalized from
Aitken et al. (1988); Bouchet et al. (1989).

##  4 Hydrogen Velocities

To understand where in the ejecta the line formation of H is occurring we
examine the profiles of all of the H lines in velocity space. Figure 8 shows
the velocity profiles of Hydrogen from 0.4-25 $\micron$. Overall, we see
evidence for $\sim$30 different spectral lines. Where four are from the
ground-based optical spectrum, three are from the ground-based NIR, 11 are
from the JWST NIRSpec data, and 11 are from the JWST MIRI spectrum. To
determine the velocity of each absorption, each H-profile was fit with a
Gaussian function using the method described in (Shahbandeh et al., 2022). The
H absorptions vary in velocity from 2100 to 4900 km s-1; see Table 4. The
velocities of the H lines are strongly dependent on the transition, whereas
the lower atomic transitions have increasingly higher velocities, see Fig. 9.
We also note that for all the H P-Cygni features, the emissions of the
profiles are blue-shifted by 1000-2000 km s-1; this is to be expected because
the extended atmosphere will block parts of the red-shifted emission.

Figure 8: The H line velocity profiles from the optical to the MIR wavelengths. We identify 29 potential H features in the ground-based optical and NIR spectra (left), the JWST NIR spectra (middle), and the JWST mid-IR spectra (Right).  Table 5: H I absorption velocities. Spectral lines affected by significant blending are indicated with an *.  Species | Rest Wavelength [$\mu$m] | Velocity [km s-1] | Error [km s-1]
---|---|---|---
Hγ | 0.434 |  $-$4900 | 900
Hβ | 0.486 |  $-$4200 | 900
Hα | 0.656 |  $-$4900 | 900
Paη* | 0.901 | 400 | 900
Paζ* | 0.923 |  $-$2300 | 900
Pa${}_{\epsilon}*$ | 0.955 |  $-$2400 | 900
Paδ | 1.005 |  $-$2700 | 200
Paγ | 1.094 |  $-$2800 | 100
Paβ | 1.282 |  $-$3100 | 100
Paα | 1.876 |  $-$4000 | 300
Brζ | 1.737 |  $-$2800 | 400
Brϵ | 1.818 |  $-$2700 | 400
Brδ | 1.945 |  $-$2800 | 400
Brγ | 2.166 |  $-$3000 | 300
Brβ | 2.626 |  $-$3300 | 300
Pfη | 2.758 |  $-$2900 | 300
Pfζ | 2.873 |  $-$2100 | 400
Pfϵ | 3.039 |  $-$2500 | 300
Pfδ | 3.297 |  $-$2700 | 400
Pfγ | 3.741 |  $-$2500 | 400
Brα | 4.052 |  $-$3600 | 400
Huη | 4.171 |  $-$2100 | 400
Huζ | 4.376 |  $-$2000 | 400
Pfβ | 4.654 |  $-$2600 | 400
Figure 9: The minimum velocity of each hydrogen feature against the transition
at which it occurs.

##  5 Full Spectral Energy Distribution

By combining all of the spectra, we can produce the first 0.4 to 25 µm SED of
a SN IIP, see Fig. 10. The average time of the SED is at $\sim 52$ d past
explosion. For orientation, we characterize the overall energy distribution by
the temperature of the black body. This is possible because Thomson and free-
free emission dominate the NIR and MIR continuum. When fit with a blackbody
function, the SED produces a temperature of $\sim$4000 K.

We compared the JWST photometry with SN 2022acko’s SED and found the NIRspec
and MIRI spectra to be consistent with the photometry, except for wavelengths
longer than 18 $\micron$, where the MIRI spectrum increases in flux. The
background in Channel 4 causes this increase in the spectral flux and is not
intrinsic to the SNe or environment. Hence, we see no evidence of pre-existing
dust or molecules in the SN 2022acko, and this statement will be quantified
below.

Figure 10: The spectral energy distribution of SN 2022acko extending between
0.4–25 $\mu$m. A constant has shifted the Keck (+ NIRES) spectrum to match the
optical NOT (+ ALFOSC) and the JWST spectra. The NIRSpec and MIRI fluxing is
consistent with the multi-band JWST photometry, except in Channel 4, where the
background dominates the spectral flux. The F2100W-band photometry sets an
important base level for future JWST observations. SN 2022acko fit with a
blackbody function characterized by $T_{eff}=4000~{}K\pm 500$ is shown in
golden.

##  6 Limits on CO formation and pre-existing molecules and dust

To obtain upper limits on the $CO$ formation, we use our framework MOLFIT
which provides an interactive layer based on the module for calculating the
vibrational and rotational bands of molecules and for spherical radiation
transport in HYDRA (see (Sharp & Höflich, 1989; Hoeflich, 1988; Gérard et al.,
2000; Höflich et al., 2002; Rho et al., 2021), and references therein). To
reproduce the observations the following free parameters are used: a)
temperature, b) density represented by a power law $\propto\rho^{n}$, c)
abundances, d) small-scale random motion such as turbulent motion
$v_{\text{turb}}$, e) average expansion velocity $v$ of a shell of thickness
$\delta v$, and f) the $CO^{+}/CO$ ratio to mimic the possible exposure to
hard $\gamma$-rays. Rather than solving the time-dependent rate equations for
molecule formation and radiation transport in HYDRA, local-thermodynamical
equilibrium is used for the molecule to atomic abundances. Note that the $CO$
concentration in LTE is an upper limit for a rapidly expanding atmospheres
because the formation times scale for $CO$ are comparable to geometrical
dilution by expansion, with an additional role of hard radiation as discussed
(see, for example Rho et al., 2021). It is well known that for the formation
and destruction, time-dependent rate equations for $C+O\rightarrow CO$,
$C^{+}+O\rightarrow CO^{+}$ are needed. The exposure of the ejecta to
$\gamma$-rays by 56Ni mixing may increase the $CO$ formation unless the
resulting destruction outweighs the formation rate via the $C+O^{+}\rightarrow
CO^{+}$ channel, and that the thermodynamic equilibrium abundance is larger
than the $CO$ abundances obtained by the time-dependent rate equations
(Spyromilio et al., 1988; Sharp & Hoeflich, 1990; Lepp et al., 1990; Meikle et
al., 1993; Rho et al., 2019). Time-dependent rate equation effects have been
included in models for SN 1987A and stripped-envelope SNe (Sharp & Höflich,
1989; Sharp & Hoeflich, 1990; Gerardy et al., 2000). Including these effects
would result in a smaller upper limit of the $CO$-mass given below. For the
simulations, 50,000 to 300,000 vibrational transitions have been taken into
account for each $CO$, $CO^{+}$ and $SiO$.

###  6.1 CO formation in the SN-ejecta

The spectrum of SN 2022acko does not show evidence for the formation of
molecules. However, our data provides important upper limits on the amount of
pre-existing molecules and dust, which is required to set a baseline for
observations at later times. In addition, it probes the mixing of C/O rich
material into the H-rich envelope. It is well-established for CCSNe with a
massive H-rich envelope, that molecule formation starts when the C/O core gets
exposed or when C/O plumes are mixed into the H-rich envelope through rayleigh
taylor instabilities. This also only occurs when the temperature drops below
$T\approx 5000$ K (Liu & Dalgarno, 1995; Gerardy et al., 2000). For example,
in SN 1987A, there was observational evidence for mixing of C/O up to
$v\approx 3000$ km s-1and H-rich material down to $v\approx 800$ km s-1(Larson
et al., 1987; Catchpole et al., 1987; Hoeflich, 1988). Mixing has been
discussed for other SNe II (Gerardy et al., 2000; Kotak et al., 2005b; Meikle
et al., 2007b; Kotak et al., 2009; Rho et al., 2018; Davis et al., 2019). This
mixing of C/O rich material was understood as due to Rayleigh-Taylor
instabilities developing in the ejecta between different chemical layers
(Muller et al., 1989; Fryxell et al., 1991; Nagataki et al., 1998).

For SN 2022acko, to obtain upper limits for $CO$ during the plateau phase, we
assumed a density structure and temperatures, $T$, typical for a SN IIP that
has an effective temperature $T\approx 5500$ K and a rapid drop in $T$ well
below 5000 K to about 2500–4000 K as the photosphere recedes (Turatto et al.,
1993; Chieffi et al., 2003; Dessart et al., 2013). For a specific example, see
Figs. 12–13 in Chieffi et al. (2003). For the density gradient, we use the
ratio of hydrogen absorption line shifts within each series of H-line features
which is limited by the optical depth of 1 in the line where each member of a
series has the same lower level or the optical depth of 1 in Thomson
scattering, as previously employed for e.g. SN 1987A (Hoeflich, 1988), and SN
1998S (Gerardy et al., 2002). Comparing the velocities obtained for
transitions with the same lower level from Table 5 and Fig. 9 and assuming
this corresponds to the point where $\tau=1$ for this transition, we obtain
$2.1<n<2.6$ assuming the density follows a powerlaw, $\rho\propto r^{-n}$
(Höflich, 1990; Duschinger et al., 1995).

During the plateau, the location of the photosphere is determined by the
recombination of H, because the opacities drop by 2 to 4 orders of magnitude
(see, for example Chieffi et al., 2003, their Figure 13).

Fig. 11 shows the scaled opacity (top) for both the fundamental (right) and
first overtone (left) of $CO$, and their flux for given temperatures
comparison with observation (bottom). The opacity peaks at about 2500–3000 K
and decreases by about six orders of magnitude towards the recombination
temperature of H. For lower temperatures, the flux in the fundamental band
scales $\propto T$ as expected, whereas the first overtone flux drops
explaining why this band vanishes at about day 200 in SNe IIP (Gerardy et al.,
2000). Because the opacity and the specific emissivity relative to the
continuum flux in the first $CO$ overtone is smaller than the fundamental band
by a factor on the order of $\approx 10^{2}$, the fundamental band provides
the upper limits, showing the importance of JWST data.

At the inner edge of the photosphere ($T\approx 5000-5500$ K), when compared
to an outer region ($T\sim 4000$ K), the $CO$ opacity is reduced by factors of
three to four orders of magnitude, thus providing a sharp inner boundary for
the $CO$ emission. For $T=4000$ K, we obtain $M(CO)\approx 2\times
10^{-6}M_{\odot}$ (Fig. 11, lower right), under the reasonable assumption that
the mixing of C/O rich material does not stop exactly at the photosphere. We
have divided our upper limit by the signal to noise ratio (S/N $\approx$44) To
obtain an upper limit for CO, we assume the potential CO feature to be
optically thin with an amplitude less than the S/N of the observed spectrum.

##  7 Comparison to SN 1987A and limits on mixing during the explosion

In SN 2022acko, the C/O-rich layers are not mixed into the H-rich layers
expanding with $\approx 2000-2500$ km s-1 as in SN 1987A. At the epoch with
similar expansion at the photosphere and light curve phase determined by the
H-recombination, SN 1987A showed extended mixing of H down to 800 km s-1 and
$CO$ out to 3000 km s-1(Danziger, 1988; Hoeflich, 1988; Meikle, 1988), whereas
SN 2022acko shows no evidence for any mixing during the explosion.

From Table 5 we see that the hydrogen lines in SN 2022acko are in a velocity
range from $\sim 2000-5000$ km s-1. The line profiles in the infared exhibit
the traditional P Cygni shape. In the case of SN 1987A the profiles, while
still generally P Cygni, showed a double-peaked emission which is very
apparent in $Pa_{\alpha}$ as well as the Balmer lines (Larson et al., 1987;
Catchpole et al., 1987). This was interpreted to be caused by mixing of C/O
into the H-rich envelope as well as of H into the C/O core (Hoeflich, 1988).
This would be from Rayleigh Taylor instabilities caused by the shock
propagation through the ejecta prior to the onset of homology (Muller et al.,
1989; Fryxell et al., 1991; Nagataki et al., 1998). Given the lack of evidence
for asymmetry in the H lines, as well as a lack of $CO$ formation, we conclude
that mixing in SN 2022acko is significantly below that of SN 1987A, likely due
to the the difference in progenitor history. Hence, any freshly formed $CO$
observable at later times in SN 2022acko should be attributed to $CO$
formation in the C/O core.

###  7.1 Molecules in the Circumstellar Material

If $CO$ exists in the CSM it would add opacity due to cold $CO$ in the
surroundings of SN 2022acko. Since we do not see $CO$, any pre-existing $CO$
is optically thin. The emission will be $\propto T$ in the fundamental band
with most of the power emitted in the first overtone of the fundamental band.
Using the scaling in Fig. 11, the upper limits of $M(CO)$ in the CSM at
$T=1000$, 500, and 100 K are $10^{-8}$, $2\times 10^{-8}$ and
$10^{-7}M_{\odot}$, respectively. If $CO$ is observed in SN 2022acko at later
times, it will be due to freshly synthesized molecules.

###  7.2 Pre-existing dust in the Interstellar and Circumstellar Medium

In principle, an upper limit for dust in the ISM and CSM can be obtained the
same way as for $CO$ above. However, the absorption to emission ratio depends
sensitively upon the grain-size distribution. For spherical dust, the surface
to mass ratio $\propto 1/R$ where $R$ is the dust radius for particles larger
than the wavelength of light or is given by Rayleigh scattering for shorter
wavelengths. In the absence of observed dust signatures, we conclude that
there is no evidence for ISM or CSM dust or a warm dust component.

Figure 11: Limits of $CO$ in SN 2022acko. $CO$ opacities cm2 g-1 Å with scale
factors of the $1^{st}$ (upper left) and fundamental band (upper right) of
$CO$, and the comparison of the continuum + $CO$ emission and observation with
the total $CO$ masses based on the fundamental band. The specific
$CO$-emissivity is given by the product of opacity and black body. For the
upper limit of the $CO$ emission, we assumed the S/N as limit for the relative
flux contribution. The underlying continuum the $CO$ feature is approximated
akin a Blackbody. The upwards slope beyond 4.8 $\mu$m hints that free-free
emission is present. Note that the mass limits based on the $1^{st}$ $CO$
would be larger by a factor of $\approx 100$, clearly demonstrating the need
for JWST observations.

##  8 Conclusions

In this work, we presented the first JWST spectrum of a CCSN, 2022acko. The
data were obtained $\sim$52 days after the explosion epoch, overlapping with
the H-recombination phase. Hydrogen lines mainly dominate the spectra of SN
2022acko but also contain spectral lines from C, N, Na, O, Mg, Br, Sc, and Fe.
Combining our JWST data with ground-based optical and NIR spectra, we created
a full SED ranging from 0.4-25 $\mu$m. The spectrum of SN 2022acko is found to
resemble similar epoch spectra of SN 1987A and SN 2005af. However, by these
phases, SN 2005af already showed forbidden lines while SN 2022acko does not.
This may suggest that SN 2022acko is evolving more slowly than SN 2005af. One
should note that the exact phase of SN 2005af is rather uncertain.

We have shown that the JWST spectra provide new constraints on the underlying
physics of SNe IIP. Our analysis of the multi-wavelength spectra allows us to
identify lines of carbon, nitrogen, and oxygen as well as s-process elements,
which confirms the expected evolutionary phases of a pre-SN red supergiant
including AGB-like burning conditions occurring in the ejecta leading to
s-process production.

We also demonstrate that placing limits on the flux level of the fundamental
band of $CO$ can be used to measure the processes happening during the stellar
evolution, the explosion, and within the environment. We see no extended
mixing of the H and the C/O core is present at this epoch, and place tight
upper limits on the $CO$ mass. We can rule out significant mixing via Raleigh
Taylor instabilities in the ejecta or a heavily asymmetric explosion. Finally,
we demonstrate that JWST is needed for finding the pre-existing $CO$ mass, and
conclude that SN 2022acko has no detectable, pre-existing $CO$
($<10^{-8}M_{\odot}$), as well as no dust in the CSM. This important finding
provides a clean slate for future JWST observations which will probe molecule
and dust formation in SN 2022acko at later times. Any molecules or dust seen
in future JWST epochs will be newly formed, since this epoch has served to
rule out the presence of any pre-existing dust or molecules.

The observations from this work set a critical baseline for molecule and dust
formation in SN 2022acko. Future JWST spectroscopic observations of SN
2022acko will provide important information on the formation of $CO$, $SiO$,
and dust. Comparison of the rate of molecule and dust formation with SNe
1987A, 2004dj, and 2005af will provide important clues into the dust budget of
SNe IIP in general as well as the variation within the SNe IIP subclass.

We thank Jozsef Vinko for providing us with electronic versions of Spitzer
spectra. MS acknowledges support by NASA grant JWST-GO-04436, JWST-GO-04217,
JWST-GO-01860, JWST-GO-02666. This work was supported by a NASA Keck PI Data
Award, administered by the NASA Exoplanet Science Institute. 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. CA, PH. EB, and JD
acknowledge support by NASA grant JWST-GO-02114, JWST-GO-02122, JWST-GO-04436,
and JWST-GO-04522. Support for program #2122 was provided by NASA through a
grant from the Space Telescope Science Institute, which is operated by the
Association of Universities for Research in Astronomy, Inc., under NASA
contract NAS 5-03127. PH acknowledges support by the National Science
Foundation (NSF) through grant AST-2306395 and NASA grant 80NSSC20K0538. EB
acknowledges support by NASA grant 80NSSC20K0538. MDS is funded by the
Independent Research Fund Denmark (IRFD, grant number 10.46540/2032-00022B).
This publication was made possible through the support of an LSSTC Catalyst
Fellowship to KAB funded through Grant 62192 from the John Templeton
Foundation to LSST Corporation. The opinions expressed in this publication are
those of the author(s) and do not necessarily reflect the views of LSSTC or
the John Templeton Foundation. ID acknowledges partial support by the Spanish
project PID2021-123110NB 100 financed by
MCIN/AEI/10.13039/501100011033/FEDER/UE. LG 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. The research of YY is supported through a Bengier-Winslow-
Robertson Fellowship. SM acknowledges support from the Magnus Ehrnrooth
Foundation and the Vilho, Yrjö and Kalle Väisälä Foundation. This work is
based on observations made with the NASA/ESA/CSA James Webb Space Telescope.
The data were obtained from the Mikulski Archive for Space Telescopes at the
Space Telescope Science Institute, which is operated by the Association of
Universities for Research in Astronomy, Inc., under NASA contract NAS 5-03127
for JWST. These observations are associated with program #2122. We acknowledge
the Keck Infrared Transient Survey (KITS), which was executed primarily by
members of the UC Santa Cruz transients team, who were supported in part by
NASA grants NNG17PX03C, 80NSSC21K2076z, 80NSSC22K1513, 80NSSC22K1518; NSF
grant AST–1911206; and by fellowships from the Alfred P. Sloan Foundation and
the David and Lucile Packard Foundation to R.J.F. KITS was directly supported
by NASA grant 80NSSC23K0301. C.D.K. is partly supported by a CIERA
postdoctoral fellowship. C.L. acknowledges support from the National Science
Foundation Graduate Research Fellowship under grant No. DGE-2233066. Y.-Z. Cai
is supported by the National Natural Science Foundation of China (NSFC, Grant
No. 12303054) and the International Centre of Supernovae, Yunnan Key
Laboratory (No. 202302AN360001).

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