# Early Spectroscopy and Dense Circumstellar Medium Interaction in SN 2023ixf

K. Azalee Bostroem LSSTC Catalyst Fellow  Steward Observatory, University of
Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  Jeniveve Pearson
Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson,
AZ 85721-0065, USA  Manisha Shrestha Steward Observatory, University of
Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  David J. Sand
Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson,
AZ 85721-0065, USA  Stefano Valenti Department of Physics and Astronomy,
University of California, Davis, 1 Shields Avenue, Davis, CA 95616-5270, USA
Saurabh W. Jha Department of Physics and Astronomy, Rutgers, the State
University of New Jersey,
136 Frelinghuysen Road, Piscataway, NJ 08854-8019, USA  Jennifer E. Andrews Gemini Observatory, 670 North A‘ohoku Place, Hilo, HI 96720-2700, USA  Nathan Smith Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  Giacomo Terreran Las Cumbres Observatory, 6740 Cortona Drive, Suite 102, Goleta, CA 93117-5575, USA  Elizabeth Green  Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  Yize Dong (董一泽) Department of Physics and Astronomy, University of California, Davis, 1 Shields Avenue, Davis, CA 95616-5270, USA  Michael Lundquist W. M. Keck Observatory, 65-1120 Māmalahoa Highway, Kamuela, HI 96743-8431, USA  Joshua Haislip Department of Physics and Astronomy, University of North Carolina, 120 East Cameron Avenue, Chapel Hill, NC 27599, USA  Emily T. Hoang Department of Physics and Astronomy, University of California, Davis, 1 Shields Avenue, Davis, CA 95616-5270, USA  Griffin Hosseinzadeh Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  Daryl Janzen Department of Physics and Engineering Physics, University of Saskatchewan, 116 Science Place, Saskatoon, SK S7N 5E2, Canada  Jacob E. Jencson Department of Physics and Astronomy, The Johns Hopkins University, 3400 North Charles Street, Baltimore, MD 21218, USA  Vladimir Kouprianov Department of Physics and Astronomy, University of North Carolina, 120 East Cameron Avenue, Chapel Hill, NC 27599, USA  Emmy Paraskeva Department of Physics and Astronomy, University of California, Davis, 1 Shields Avenue, Davis, CA 95616-5270, USA  Nicolas E. Meza Retamal Department of Physics and Astronomy, University of California, Davis, 1 Shields Avenue, Davis, CA 95616-5270, USA  Daniel E. Reichart Department of Physics and Astronomy, University of North Carolina, 120 East Cameron Avenue, Chapel Hill, NC 27599, USA  Iair Arcavi School of Physics and Astronomy, Tel Aviv University, Tel Aviv 69978, Israel  Alceste Z. Bonanos IAASARS, National Observatory of Athens, Penteli 15236, Greece  Michael W. Coughlin School of Physics and Astronomy, University of Minnesota, 116 Church Street S.E., Minneapolis, MN 55455, USA  Ross Dobson Mullard Space Science Laboratory, University College London, Holmbury St Mary, Dorking, Surrey RH5 6NT, United Kingdom  Isaac Newton Group of Telescopes, Apt. de Correos 368, E-38700 Santa Cruz de la Palma, Spain  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  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, Gran Capità, 2-4, Edifici Nexus, Desp. 201, E-08034 Barcelona, Spain  Claudia Gutiérrez Institut d’Estudis Espacials de Catalunya, Gran Capità, 2-4, Edifici Nexus, Desp. 201, E-08034 Barcelona, Spain  Institute of Space Sciences (ICE, CSIC), Campus UAB, Carrer de Can Magrans, s/n, E-08193 Barcelona, Spain  Suzanne Hawley  Department of Astronomy, University of Washington, 3910 15th Avenue NE, Seattle, WA 98195-0002, USA  Leslie Hebb Physics Department, Hobart and William Smith Colleges, 300 Pulteney Street, Geneva, NY 14456, USA  Department of Astronomy, Cornell University, 245 East Avenue, Ithaca, NY 14850, USA  Daichi Hiramatsu Center for Astrophysics | Harvard & Smithsonian, 60 Garden Street, Cambridge, MA 02138-1516, USA  The NSF AI Institute for Artificial Intelligence and Fundamental Interactions, 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  Takashi Iijima  INAF-Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5, I-35122 Padova, Italy  Ilya Ilyin Leibniz-Institut für Astrophysik Potsdam (AIP), An der Sternwarte 16, D-14482 Potsdam, Germany  Kiran Jhass Department of Physics and Astronomy, University of Sheffield, Sheffield S3 7RH, UK  Isaac Newton Group of Telescopes, Apt. de Correos 368, E-38700 Santa Cruz de la Palma, Spain  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  Sean Moran Center for Astrophysics | Harvard & Smithsonian, 60 Garden Street, Cambridge, MA 02138-1516, USA  Brett M. Morris Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218-2410, USA  Alessandra C. Mura Physics and Astronomy Department Galileo Galilei, University of Padova, Vicolo dell’Osservatorio 3, I-35122, Padova, Italy  Tomás E. Müller-Bravo Institute of Space Sciences (ICE, CSIC), Campus UAB, Carrer de Can Magrans, s/n, E-08193 Barcelona, Spain  Institut d’Estudis Espacials de Catalunya, Gran Capità, 2-4, Edifici Nexus, Desp. 201, E-08034 Barcelona, Spain  James Munday Department of Physics, Gibbet Hill Road, University of Warwick, Coventry CV4 7AL, United Kingdom  Isaac Newton Group of Telescopes, Apt. de Correos 368, E-38700 Santa Cruz de la Palma, Spain  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  Maria Th. Pabst Physics and Astronomy Department Galileo Galilei, University of Padova, Vicolo dell’Osservatorio 3, I-35122, Padova, Italy  Paolo Ochner INAF-Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5, I-35122 Padova, Italy  Physics and Astronomy Department Galileo Galilei, University of Padova, Vicolo dell’Osservatorio 3, I-35122, Padova, Italy  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  Andrea Pastorello INAF-Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5, I-35122 Padova, Italy  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  Lara Piscarreta Institute of Space Sciences (ICE, CSIC), Campus UAB, Carrer de Can Magrans, s/n, E-08193 Barcelona, Spain  Aravind P. Ravi Department of Physics, University of Texas at Arlington, Box 19059, Arlington, TX 76019, USA  Andrea Reguitti INAF-Osservatorio Astronomico di Brera, Via E. Bianchi 46, I-23807, Merate (LC), Italy  INAF-Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5, I-35122 Padova, Italy  Laura Salo School of Physics and Astronomy, University of Minnesota, 116 Church Street S.E., Minneapolis, MN 55455, USA  József Vinkó University of Texas at Austin, 1 University Station C1400, Austin, TX 78712-0259, USA  Konkoly Observatory, CSFK, MTA Center of Excellence, Konkoly-Thege M. út 15-17, Budapest, 1121, Hungary  ELTE Eötvös Loránd University, Institute of Physics and Astronomy, Pázmány Péter sétány 1/A, Budapest, 1117 Hungary  Department of Experimental Physics, University of Szeged, Dóm tér 9, Szeged, 6720, Hungary  Kellie de Vos School of Physics and Astronomy, University of Nottingham, University Park, Nottingham, NG7 2RD  Isaac Newton Group of Telescopes, Apt. de Correos 368, E-38700 Santa Cruz de la Palma, Spain  J. C. Wheeler University of Texas at Austin, 1 University Station C1400, Austin, TX 78712-0259, USA  G. Grant Williams Steward Observatory, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  MMT and Steward Observatories, University of Arizona, 933 North Cherry Avenue, Tucson, AZ 85721-0065, USA  Samuel Wyatt Department of Astronomy, University of Washington, 3910 15th Avenue NE, Seattle, WA 98195-0002, USA

###### Abstract

We present the optical spectroscopic evolution of SN 2023ixf seen in sub-night
cadence spectra from 1.18 to 15 days after explosion. We identify high-
ionization emission features, signatures of interaction with material
surrounding the progenitor star, that fade over the first 7 days, with rapid
evolution between spectra observed within the same night. We compare the
emission lines present and their relative strength to those of other
supernovae with early interaction, finding a close match to SN 2020pni and SN
2017ahn in the first spectrum and SN 2014G at later epochs. To physically
interpret our observations we compare them to CMFGEN models with confined,
dense circumstellar material around a red supergiant progenitor from the
literature. We find that very few models reproduce the blended N III
($\lambda\lambda 4634.0,4640.6$)/C III ($\lambda\lambda 4647.5,4650.0$)
emission lines observed in the first few spectra and their rapid disappearance
thereafter, making this a unique diagnostic. From the best models, we find a
mass-loss rate of $10^{-3}-10^{-2}$ $\mathrm{M_{\odot}\,yr^{-1}}$, which far
exceeds the mass-loss rate for any steady wind, especially for a red
supergiant in the initial mass range of the detected progenitor. These mass-
loss rates are, however, similar to rates inferred for other supernovae with
early circumstellar interaction. Using the phase when the narrow emission
features disappear, we calculate an outer dense radius of circumstellar
material $R_{\mathrm{CSM,out}}\approx 5\times 10^{14}~{}\mathrm{cm}$ and a
mean circumstellar material density of $\rho=5.6\times
10^{-14}~{}\mathrm{g\,cm^{-3}}$. This is consistent with the lower limit on
the outer radius of the circumstellar material we calculate from the peak
$\mathrm{H\alpha}$ emission flux, $R_{\text{CSM, out}}\gtrsim 9\times
10^{13}~{}\mathrm{cm}$.

Core-collapse supernovae (304), Type II supernovae (1731), Circumstellar
matter (241), Stellar mass loss (1613), Red supergiant stars (1375)

††facilities: ADS, ARC (ARCES), Asiago:Galileo (B&C), Bok (B&C), HCT (HFOSC),
HET (LRS2), INT (IDS), LBT (MODS), LCOGT (FLOYDS), Liverpool:2m (SPRAT), MMT
(Hectospec), NED, NOT (ALFOSC), TNS††software:  aesop (Morris & Dorn-
Wallenstein, 2018), astropy (; ; ), CMFGEN (Hillier & Miller, 1998; Hillier &
Dessart, 2012, 2019; Dessart et al., 2013), FLOYDS (Valenti et al., 2014),
HSRED, IRAF (Tody, 1986, 1993), Light Curve Fitting (Hosseinzadeh et al.,
2023), MatPLOTLIB (Hunter, 2007), MODS pipeline (Pogge, 2019), NumPy (Harris
et al., 2020), Scipy (Virtanen et al., 2020).

##  1 Introduction

Type II supernovae (SNe; hydrogen-rich, specifically Type IIP/L) are thought
to come from red supergiant (RSG) progenitors, with masses of $\sim$8–25 M⊙
(Smartt et al., 2009; Smartt, 2015). While there is a consensus that massive
stars enrich their environments through mass loss, there is no model that
quantitatively predicts observed RSG mass-loss rates (Kee et al., 2021, and
references therein). Empirical mass-loss rates derived from direct
observations span orders of magnitude (Mauron & Josselin, 2011). While the
empirical prescription most often used in single-star evolutionary models is
that of de Jager et al. (1988) (which spans 10-7–10-3.8
$\mathrm{M_{\odot}\,yr^{-1}}$ for RSG luminosities of $\log(L/L_{\odot})=$
3.9–5.8) recent analyses have found evidence of significantly higher (Ekström
et al., 2012; Massey et al., 2023) and lower (Beasor et al., 2020) mass-loss
rates. Recently, observations of both the early light curves and spectra of
Type II SNe show evidence of dense circumstellar material (CSM; Khazov et al.,
2016; Morozova et al., 2018; Bruch et al., 2022; Subrayan et al., 2023),
indicating more extreme mass loss (e.g. eruptive mass loss or a superwind),
which are known to occur in more massive stars (e.g. luminous blue variables;
LBVs Smith et al., 2011). Regardless of the mechanism, all massive stars lose
mass, and therefore we expect the photons and ejecta from all of their
resultant SN explosions to interact with the circumstellar material (CSM)
surrounding the progenitor at some level (see Smith, 2014, for a review).

One of the signatures of CSM interaction in core-collapse (CC) SNe is narrow
emission lines corresponding to highly-ionized species in their early spectra.
Narrow emission lines can first occur when photons from shock breakout ionize
surrounding CSM (e.g. Yaron et al., 2017). As the shock passes through the
CSM, the kinetic energy of the ejecta is converted to high-energy photons that
can also ionize the CSM ahead of the photosphere (e.g. Leonard et al., 2000;
Smith et al., 2015; Terreran et al., 2022). The recombination of the ionized
gas leads to emission features, with the photons scattering off of the ionized
electrons to produce Lorentzian line wings (Chugai, 2001). This produces high-
ionization features which are a function of the temperature, density, and
composition of the CSM (e.g. O VI, O V, N V, N IV, C IV, He II). As the CSM
cools, higher ionization features give way to lower ionization features (e.g.
N III, O III) and eventually all emission lines fade. At the same time, narrow
P Cygni profiles can develop if the CSM is dense enough and sufficiently cool
(e.g. Leonard et al. 2000; Benetti et al. 2016; Terreran et al. 2022). These
profiles can develop into intermediate width features as the CSM begins to be
accelerated by the shock. Eventually, the ejecta sweep up or engulf the CSM
and the spectrum begins to develop as a normal CCSN with broad P Cygni
profiles, often with a shallow absorption component. However, even at this
phase the CSM interaction can contribute to the light curve (Dessart &
Hillier, 2022; Andrews & Smith, 2018; Smith et al., 2015; Smith, 2017).

Analyses of samples of CCSNe with early spectroscopic observations show that a
significant fraction of nearby SNe display these features (Khazov et al.,
2016; Bruch et al., 2021, 2022). Detailed modeling of these flash features can
constrain the progenitor mass-loss rate just prior to explosion, the surface
chemical composition, as well as the extent of the confined CSM (e.g. Dessart
et al., 2017; Boian & Groh, 2019, 2020). There are now dozens of examples of
early flash spectroscopy, and several cases where the observations have been
modeled or been compared to models in some detail (e.g. Yaron et al., 2017;
Boian & Groh, 2020; Tartaglia et al., 2021; Terreran et al., 2022; Jacobson-
Galán et al., 2022), but the time evolution of the flash ionization lines has
rarely been captured due to their ephemeral nature.

SN 2023ixf was discovered in M101 (D=6.85 Mpc, Riess et al., 2022) by Koichi
Itakagi on 2023-05-19 17:27:15 UTC (all times given in this paper are in UTC;
MJD 60083.72) at a magnitude of 14.9 AB mag in a Clear filter (Itagaki, 2023).
It was classified on 2023-05-19 23:35:34 (MJD 60083.98) as a Type II SN with
flash ionization features (H, He, C, and N), using a spectrum taken a few
hours after discovery (Perley & Gal-Yam, 2023). Over the first $\sim$5 days,
it rapidly rose to a plateau brightness of $V\approx$ 11.2 mag, or
$M_{V}\approx-18.2$ at the distance to M101 – a similar brightness to the
well-studied Type IIP SN 2004et in NGC 6946 (e.g. Maguire et al., 2010). The
early photometric evolution is detailed in a companion paper, by Hosseinzadeh
et al. (2023).

In this paper, we present the remarkable early spectroscopic evolution of SN
2023ixf with flash features observed in extraordinary detail. In Section 2, we
describe our spectroscopic observations, while in Section 3 we present basic
properties of SN 2023ixf relevant for our work. From there, in Section 4 we
discuss the fast spectroscopic evolution of SN 2023ixf and compare it with
existing observational data sets. In Section 5, we compare our unprecedented
flash spectroscopic sequence to existing radiative transfer models to infer
the mass-loss rate of the progenitor star and in Section 6, we use the
spectroscopic evolution to further characterize the CSM of SN 2023ixf and
place it in the context of other interacting SNe. We summarize and conclude in
Section 7.

##  2 Spectroscopic Observations

Immediately after the discovery announcement, we began a high-cadence,
comprehensive campaign to observe the detailed evolution of SN 2023ixf with
the Arizona Transient Exploration and Characterization (AZTEC) collaboration,
Distance Less Than 40 Mpc (DLT40) collaboration, and the Global Supernova
Project. We observed SN 2023ixf using the moderate-resolution optical
spectrograph Hectospec (Fabricant et al., 2005) on the Multiple Mirror
Telescope (MMT) on Mt. Hopkins, AZ from 2023-05-20 to 2023-05-26. The observed
spectra were reduced using an IDL pipeline called
HSRED111http://mingus.as.arizona.edu/~bjw/mmt/hecto_reduction.html and then
flux calibrated using IRAF (Tody, 1986, 1993). Further optical spectroscopy
was obtained using the FLOYDS spectrograph on Faulkes Telescope North (FTN)
through the Global Supernova Project. Spectra were reduced with standard
methods using a custom IRAF-based pipeline (Valenti et al., 2014). Adding to
our high-cadence spectroscopic coverage, we observed SN 2023ixf with the
Alhambra Faint Object Spectrograph (ALFOSC) on the Nordic Optical Telescope
(NOT; Proposal 67-112, PI: Bonanos). Observations were reduced with standard
reduction techniques using IRAF. We add publicly available ALFOSC observations
from the NUTS2 collaboration222The Nordic optical telescope Unbiases
Transients Surveys 2; https://nuts.sn.ie (Stritzinger et al., 2023). In
addition, we observed SN 2023ixf with the Multi-Object Double Spectrographs
(Pogge et al., 2010) on the Large Binocular Telescope (LBT). Data were bias
and flat-field corrected using the modsCCDred package (Pogge, 2019), then
extracted and flux calibrated with IRAF. Spectroscopic observations were taken
with the Boller and Chivens Spectrograph (B&C) on University of Arizona’s Bok
2.3m telescope located at Kitt Peak Observatory. These observations were
reduced using standard IRAF reduction techniques. An optical spectrum was also
taken with the Low Resolution Spectrograph 2 (LRS2, Chonis et al., 2016) on
the Hobby-Eberly Telescope (HET, Ramsey et al., 1998; Hill et al., 2021) at
McDonald Observatory on 2023-06-02. The data from the red and blue arms
(LRS2-R and LRS2-B) were combined into a single spectrum covering the spectral
region from 3600 to 10500 Å. The Integral Field Unit (IFU) spectra were
reduced with the Panacea pipeline333https://github.com/grzeimann/Panacea. We
collected one 30-minute exposure with the Astrophysical Research Consortium
Echelle Spectrograph (ARCES) with resolution $R\approx 31,000$ on the ARC 3.5
m Telescope at Apache Point Observatory (APO). We reduced the spectra with
IRAF and aesop (Morris & Dorn-Wallenstein, 2018). Individual orders normalized
using a Spline function. Further optical spectra were obtained with the 1.22-m
Galileo telescope+B&C at the Asiago Astrophysical Observatory, Italy, which
were reduced using an IRAF-based pipeline and the Intermediate Dispersion
Spectrograph (IDS) at the Isaac Newton Telescope (INT) which were reduced with
the custom python package IDSRED (Bravo, 2023). . To this dataset we add
publicly available reduced spectroscopic observations from the Liverpool
Telescope (LT) archive (SPRAT; Perley, 2023; Steele et al., 2004; Piascik et
al., 2014) and Transient Name Server444https://www.wis-tns.org/object/2023ixf
(HFOSC, SPRAT).

A complete list of spectroscopic observations is given in Appendix A and shown
in Figure 1.1 and Figure 1.2. Spectra will be made available on
Wiserep555https://www.wiserep.org (Yaron & Gal-Yam, 2012).

Figure 1.1: The evolution of the optical spectra of SN 2023ixf over the first
$\sim 5.2$ days, with the earliest epoch at the top of the figure. Spectra are
color coded by the instrument used to observe them, normalized by a black body
fit, and corrected for redshift. Emission lines are identified at their rest
wavelengths with vertical lines and labeled at the top of the figure, while
the most prominent telluric features are marked with the shaded gray region.
Throughout this sequence, the spectra evolve from showing strong, narrow
emission lines from high-ionization species to a intermediate width features
as the CSM is accelerated by the shock. The spectra evolve rapidly between
1.18–1.67d, with the He I and N III ($\lambda\lambda 4634.0,4640.6$)/C III
($\lambda\lambda 4647.5,4650.0$) lines disappearing within the first 0.5d from
the first spectrum while He II ($\lambda 4685.5$), N IV, and C IV gain
strength.  Figure 1.2: Continuation of Figure 1.1 showing the evolution of the
optical spectra of SN 2023ixf from day 5.2 to 14.5 with the earliest epoch at
the top of the figure. Over the first 7 days, the spectra evolve from showing
strong, narrow emission lines from high-ionization species to a nearly
featureless spectrum with an intermediate with P Cygni profile in
$\mathrm{H\alpha}$. In the subsequent 7 days, broad P Cygni profiles develop
in the higher order Balmer features.

##  3 Fundamental Supernova Parameters

The nearby host galaxy of SN 2023ixf, M101, also hosted the Type Ia SN 2011fe.
We adopt the distance modulus to M101, derived using the Leavitt Law, from
Riess et al. (2022): $\mu=29.178\pm 0.041$ mag (D=6.85 Mpc). For Milky Way
extinction, we use the SN location and the dust maps of Schlafly & Finkbeiner
(2011)666https://irsa.ipac.caltech.edu/applications/DUST/index.html to find
$E(B-V)=0.0077\pm 0.0002$ mag. Measuring the equivalent width of the Na I D
lines in high-resolution observations and using the relationship of Poznanski
et al. (2012), Smith et al. (2023) find an average host $E(B-V)=0.031$ mag
with $\pm$30% uncertainty from the relation between Na i D and extinction; a
value which is consistent with other Na I D extinction measurements made with
high resolution data (Lundquist et al., 2023).

M101 is a popular target for both amateur and professional observers, and
images of the galaxy taken by amateur astronomers (Mao et al., 2023) prior to
discovery provided the last deep non-detection (2023-05-18 15:50:24; MJD
60082.66) and first detection (2023-05-18 20:29; MJD 60082.85), with
$\lesssim$5 hours separating them. Following Hosseinzadeh et al. (2023), we
define the explosion epoch as half way between the last non-detection and
first detection: 2023-05-18 18:00:00 (MJD 60082.75 $\pm$ 0.10), where the
adopted uncertainty is the span between the explosion epoch and last non-
detection (or first detection).

##  4 Spectroscopic Evolution

Tracking the rapid evolution of SN 2023ixf, we observed SN 2023ixf at least
four times per night for the first 5 days and at least nightly thereafter.
While high-cadence spectra have been obtained for a select few SNe over the
first few days of evolution (Yaron et al., 2017; Terreran et al., 2022), SN
2023ixf is the first to have intra-night observations for the first week. Over
the first two weeks, the spectra evolve from strong, narrow emission lines
with broad wings, to a nearly featureless spectrum and finally develop Balmer
P Cygni profiles with shallow absorption, more typical of the early evolution
of Type II SNe.

To identify spectral features, we use the second spectrum, taken 1.36 days
after explosion, as it has higher signal-to-noise and resolution than the
classification spectrum. From this spectrum, we identify H I, He I, He II, C
III, C IV, O III, N III, and N IV lines. A full list of species is given in
Appendix B.

We see a rapid evolution in the first 0.5 days of our spectra (1.18–1.67
days), which is shown in detail in Figure 2. First we turn to the complex of
lines around 4700 Å in the top panel of the figure. Over this epoch, the blend
of N III ($\lambda\lambda 4634.0,4640.6$)/C III ($\lambda\lambda
4647.5,4650.0$) fades into the broad blueshifted wing of He II ($\lambda
4685.5$) while the latter line increases in strength. Similarly, the He I
lines visible in the first spectrum rapidly fade until they are no longer
visible 0.5d later. Other lines such as O III, C III, and N IV ($\lambda$ 5074
Å) that are marginally detected in the first few spectra also disappear on
this time scale. At the same time, N IV ($\lambda\lambda 7103,7109,\lambda
7122$) along with He II ($\lambda 4685.5$), C IV ($\lambda\lambda$ 5801.3,
5811.98) increase in strength.

The most persistent features through day 5 are H I ($\lambda$ 4101.73) and
possibly N III ($\lambda$ 4097.33; although it is not possible to determine if
two distinct lines exist at this resolution), He II ($\lambda 4685.5$), and C
IV ($\lambda\lambda 5801.3,5811.98$). Interestingly, N III ($\lambda\lambda
4634.0,4640.6$)/C III ($\lambda\lambda 4647.5,4650.0$) become visible again as
a broad shelf in the He II ($\lambda 4685.5$) blue wing around day 4, as the
He II ($\lambda 4685.5$) line fades.

Over this time, the strength of all these narrow features decreases to a
nearly featureless spectrum 7 days after explosion. The one line still clearly
present in the spectrum after day 6.5 is $\mathrm{H\alpha}$, which appears to
develop an asymmetric emission profile, reminiscent of a P Cygni profile with
shallow absorption. The intermediate width of this feature indicates the
presence of CSM that has been accelerated by the shock. Details of its
evolution are presented by Smith et al. (2023). Around day 12, clear broad P
Cygni profiles from the SN ejecta are present in the high-order Balmer
features, although the profile at $\mathrm{H\alpha}$ is significantly more
complex with no clear absorption component.

Figure 2: The evolution of SN 2023ixf from day 1.18 through day 3, colored by
phase. Ions are labeled in each panel and the phase is given in the middle
panel. Emission lines from the low-ionization levels, He I, N III, C III, O
III, and N IV ($\lambda$ 5074 Å), disappear over the first 0.5 days of
evolution while emission from high-ionization levels, N IV ($\lambda\lambda
7103,7109,\lambda 7122$), He II ($\lambda 4685.5$), and C IV ($\lambda\lambda$
5801.3, 5811.98 Å) increase in strength. Spectra are fit with a blackbody
function and normalized, corrected for redshift, and offset for readability.
In the bottom panel, the B-band telluric feature is marked with a gray shaded
region.

###  4.1 Comparison to other Flash SNe

A number of SNe have been observed to have narrow emission lines early in
their evolution, often disappearing within the first week. In Figure 3 we
compare the spectra of SN 2023ixf at day 1, 3, 7, and 14 to those of SN 1998S
(Leonard et al., 2000), SN 2013fs (Yaron et al., 2017), SN 2013cu (Gal-Yam et
al., 2014), SN 2014G (Terreran et al., 2016), SN 2017ahn (Tartaglia et al.,
2021), and SN 2020pni (Terreran et al., 2022) at comparable phases.

Given the rapid evolution within the first week of explosion, the uncertainty
in the explosion epoch of each comparison supernova should be considered. SN
1998S was discovered on 1998-03-02 16:19:12 with the last non-detection 7 days
earlier on 1998-02-23 16:47:24. With the discovery image taken 18 days before
V band maximum, Leonard et al. (2000), conclude that it was likely discovered
within a few days of shock breakout. We take the explosion epoch to be half
way between the last non-detection and first detection (1998-02-27 04:33:18)
with an uncertainty of the distance to the last non-detection/first detection:
3.5d. SN 2013cu was discovered on 2013-05-03 04:19:12. Gal-Yam et al. (2014)
derive an explosion epoch of 2013-05-02 22:19:12 with an uncertainty of 0.11d
which we adopt in this paper. Similarly, Yaron et al. (2017) find an explosion
epoch of 2013-10-06 02:52:48, a few hours prior to the first detection on
2013-10-06 06:05:45 with an uncertainty of 0.5 hr which we use in this
analysis. Terreran et al. (2016) set the explosion epoch of SN 2014G to
2014-01-12 14:24:00, half way between the last non-detection and first
detection, with an uncertainty of 1.7d. SN 2017ahn was discovered on
2017-02-08 06:57:00, 1 day after the last non-detection on 2017-02-07
05:31:12. Tartaglia et al. (2021) find the explosion epoch to be half way
between the last non-detection and first detection: 2017-02-07 18:14:24 with
an uncertainty of 0.5d. For SN 2020pni, Terreran et al. (2022) determine an
explosion epoch of 2020-07-15 19:12:00, 0.4d prior to the last non-detection,
with an uncertainty of 0.1d.

These SNe encompass a range of subtypes of Type II supernovae which are
defined by their light curve shape or spectral characteristics. Type IIP/L SNe
are defined by hydrogen in their spectra throughout their evolution and a
plateau from peak brightness to $\sim$100 days after explosion which is either
flat or linearly declining, with a continuum of possible slopes. One
explanation for the different slopes is that steeper slopes represent
progenitor stars with smaller hydrogen envelopes due to mass loss (Grassberg
et al., 1971; Young & Branch, 1989; Blinnikov & Bartunov, 1993; Anderson et
al., 2014; Moriya et al., 2016; Hiramatsu et al., 2021). In the cases where
only a very small hydrogen envelope remains, the supernova becomes a Type IIb
SN, defined by the transition of their spectra from hydrogen dominated to
showing strong helium lines (Filippenko, 1997). All of the comparison SNe are
Type IIP/L SNe with the exception of SN 2013cu, which is a Type IIb (Leonard
et al., 2000; Gal-Yam et al., 2014; Terreran et al., 2016; Yaron et al., 2017;
Tartaglia et al., 2021; Terreran et al., 2022). Their light curves show
similarly shallow linear declines in V-band, with SN 2013cu being the steepest
(Yamanaka et al., 2023; Teja et al., 2023; Jacobson-Galan et al., 2023;
Hosseinzadeh et al., 2023; Hiramatsu et al., 2023; G. et al., 2023).

Among the SNe in this sample, SN 2013fs is the most distinct. It showed the
highest ionization species (e.g. O V; $\lambda 5541$ and O VI; $\lambda\lambda
3811,3834$; Yaron et al. 2017) in the earliest spectra that are not present in
SN 2023ixf or other SNe in this sample. However, it is possible that these
features were present at earlier times in SN 2023ixf and, by the first
spectrum at day 1.18, had disappeared. While there are emission lines from
some of the same ions in the day 1.5 spectra of SN 2013fs and SN 2023ixf, many
are not present in the SN 2013fs spectrum, especially between 3800–4500 Å and
5000–6500 Å, and N V ($\lambda$ 4604) is absent from the SN 2023ixf spectrum.
Additionally, the emission in all lines is weaker in SN 2013fs and fades more
rapidly.

SN 2023ixf more closely resembles the other comparison objects in Figure 3,
although it does not always match their evolution. The first spectrum closely
resembles that of SNe 2020pni and 2017ahn. However, some features, e.g. N III
($\lambda\lambda 4634.0,4640.6$)/C III ($\lambda\lambda 4647.5,4650.0$), are
persistent at day 3 in SNe 2020pni and 2017ahn, but are not visible in SN
2023ixf. This likely indicates that the CSM of SN 2023ixf is cooling more
slowly than that of SN 2020pni or SN 2017ahn, perhaps due to a lower density.
At day 3, we also have spectra from SN 2013cu and SN 2014G. At this phase, SN
2023ixf is best matched to SN 2014G, a trend that continues throughout the two
week evolution, while SN 1998S and SN 2013cu more closely resemble SN 2020pni
and SN 2017ahn. At day 7, most features have faded and $\mathrm{H\alpha}$ has
been replaced by an intermediate width P Cygni profile in SN 2023ixf, in
contrast to the narrow P Cygni profiles in SNe 2020pni and 1998S. Finally, at
day 14, SN 2020pni and SN 1998S still show narrow P Cygni profiles, while SN
2023ixf and SN 2014G are starting to develop broad P Cygni profiles. The
disappearance of these features can be interpreted as the ejecta enveloping
the CSM, which would indicate that the radial extent of the CSM of SN 2023ixf
is smaller than that of SN 2020pni and SN 1998S or the ejecta velocity higher.

Figure 3: A comparison of SN 2023ixf (black) to SN 1998S (teal; Leonard et
al., 2000), SN 2013fs (blue; Yaron et al., 2017), SN 2013cu (green; Gal-Yam et
al., 2014), SN 2014G (purple; Terreran et al., 2016), SN 2017ahn (mustard;
Tartaglia et al., 2021), and SN 2020pni (pink; Terreran et al., 2022) $\sim$1,
3, 7, and 14 days after explosion. These SNe all have flash features that
evolve over the first two weeks. The SN name and epoch are marked to the right
of the plot. All spectra have been redshift and extinction corrected and then
normalized by a blackbody fit to the continuum. Emission lines identified in
the first spectrum are marked with vertical lines which are labeled at the top
of the figure. In the first epoch, SN 2023ixf closely resembles SNe 2020pni
and 2017ahn, although at later epochs it more closely follows SN 2014G. SN
2013fs is considerably different in its evolution throughout the first two
weeks. At later epochs, SN 2020pni and SN 1998S develop narrow P Cygni
profiles that are not present in the spectra of other SNe.

##  5 Comparison to Models

Variations in mass-loss rate, SN luminosity, surface abundance, and CSM
density profile can all affect the characteristics of the narrow emission line
spectrum and its evolution. In the following section, we compare our
spectroscopic dataset to two sets of publicly available model grids (Boian &
Groh, 2020; Dessart et al., 2017), which vary different parameters, using the
best-fit model to characterize the progenitor, CSM, and SN properties. We fit
a blackbody to both the model and observed data and normalize by this before
comparing them. This removes any temperature continuum effects, instead only
examining flux relative to the continuum.

###  5.1 Boian & Groh models

Boian & Groh (2019) propose that the narrow features produced by the CSM
around Type II SNe can be used to constrain the abundances of this material
and therefore the mass of the progenitor system. They predict different SN
line diagnostics for low-mass RSG progenitors (8–15 M⊙), massive RSG, yellow
hypergiant, blue supergiant progenitors (15–30 M⊙), and stripped stars like
LBV and N-rich Wolf-Rayet progenitors (15–30 M⊙) for observations from one to
a few days after explosion. These progenitors are then modeled for high-,
medium-, and low-luminosity SNe using the radiation transport code CMFGEN
(Hillier & Miller, 1998; Hillier & Dessart, 2012; Dessart et al., 2013;
Hillier & Dessart, 2019). The primary difference between the low- and high-
mass progenitors is in the abundances of the surface material. Low-mass
progenitors should experience weak or no CNO processing. High mass
progenitors, on the other hand, are expected to have significant CNO processed
materials. Finally, stripped stars should be He-rich. Table 1 gives the
predicted spectroscopic signatures of low and high-mass progenitors and
different luminosity SNe.

We search for these features in our 1.36d spectrum. We see no O VI
($\lambda\lambda 3811,3834$) which rules out low-mass, high-luminosity SNe as
this is the only diagnostic predicted for this combined mass and luminosity.
We also do not see C III ($\lambda 5697$), which is one of three diagnostics
of the low-mass, medium-luminosity system. It is unclear how to interpret a
scenario in which only a subset of the diagnostic lines are present.
Additionally, we identify N IV ($\lambda 4058$, $\lambda\lambda 7109,7122$), N
III ($\lambda\lambda 4634,4640$), C III ($\lambda\lambda 4647,4650$), and C IV
($\lambda\lambda$ 5801, 5811), which are features of the remaining progenitor
mass - luminosity combinations. Thus, although we can eliminate a low mass,
high luminosity progenitor, our conclusions about the remaining combinations
of mass and luminosity are inconclusive as the diagnostic lines for multiple
scenarios are observed and for some of them only a subset of the features are
present. We note that this diagnostic depends on uncertain physical parameters
in single star evolutionary models such as mixing efficiency, mass-loss rates,
and convective overshoot which complicate the connection between surface
abundance, as measured from CSM, and initial mass. Additionally, the
temperature of the CSM has a dramatic effect on the ions present, independent
of surface abundances, which is a function of both the individual supernova
parameters as well as the phase at which the model and observations are
compared. This further complicates the connection between surface abundances
and progenitor masses.

As the broad trends noted by Boian & Groh (2019) were inconclusive, we
visually compare their full set of model spectra for all mass-loss rates,
luminosities, and progenitors to our 1.36d spectrum. The clear presence of
both N III ($\lambda$4634) and He II ($\lambda$4685.5) with He II stronger
than N III greatly limits the number of possible models to $L=1.5\times
10^{9}$ $\mathrm{L_{\odot}}$ (which they define as medium luminosity) and
$\dot{M}=3\times 10^{-3}$ M⊙ yr-1 (these lines are not present in any other
models). The wind velocity of these models is
$v_{w}=150~{}\mathrm{km\,s^{-1}}$ which is consistent with the observed CSM
velocity (Smith et al., 2023).

While we have identified a mass-loss rate and luminosity which is consistent
with our observed spectrum based on the presence of individual features, Boian
& Groh (2020) scale the model luminosity to match the observed luminosity,
which leads to a scaling in the mass-loss rate as well. We repeat this
analysis to derive a luminosity and mass-loss rate for SN 2023ixf. First,
Boian & Groh (2020) identify the temperature via the ionization level present.
Similarly, we selected the best models based on continuum normalized spectra,
looking only at the relative ionization levels. In their model, the
temperature and luminosity are related via the Stefan-Boltzmann law. Thus
after identifying the temperature, they scale to the observed luminosity,
scaling the radius to maintain the derived temperature. They also derive a
scale factor for the mass-loss rate, assuming that $\dot{M}\propto
L_{SN}^{3/4}$. We duplicate this analysis for SN 2023ixf for the three surface
abundances (which do not produce significantly different results), concluding
that the mass-loss rate of the progenitor of SN 2023ixf was $\dot{M}\approx
4.5\times 10^{-3}$ $\mathrm{M_{\odot}\,yr^{-1}}$ and scaled luminosity for
each abundance is $L\approx 2.6\times 10^{9}$ $\mathrm{L_{\odot}}$.

Having constrained the luminosity and mass-loss rate, we examine the different
surface abundances: the solar abundance spectra, corresponding to the low mass
RSG scenario; the CNO-processed surface abundance, corresponding to the high-
mass RSG, BSG, YSG scenario; and the He-rich abundance, corresponding to the
LBV, WN, stripped star scenario. These are compared to an observed spectrum in
Figure 4. We find that no one scenario matches the line strengths of the
observation. While many of the lines are well matched in the solar abundance
model, the C IV ($\lambda\lambda 5801,5811$) and $\mathrm{H\alpha}$ are
greatly over-estimated by the model. The C IV ($\lambda\lambda 5801,5811$) is
better represented in the CNO abundance model (a result of the suppression of
C and O), however, the N in our observed spectrum is significantly weaker than
the model, countering the expectation that these stars would be nitrogen
enriched. Finally, in the He-rich model, the H I line is well-modeled and the
He II ($\lambda 4685.5$) line is stronger in the model than in the
observation. Like the CNO abundance model, the N lines are too strong in the
model. It is possible that, rather than indicating surface abundance, these
discrepancies arise from a mismatch between the physical and model CSM
density, temperature, and/or the hardness of the radiation field.

Table 1: Predicted spectroscopic signatures from Boian & Groh (2019).  Luminosity | Low-mass Progenitor (8–15 M⊙) | High-mass Progenitor (15–30 M⊙)
---|---|---
Low luminosity | C III ($\lambda\lambda 4647,4650$) | C IV ($\lambda\lambda 5801,5811$)
($\sim 1.9\times 10^{9}$ $\mathrm{L_{\odot}}$) |  | Lack of C III ($\lambda 5697$)
Medium luminosity | C III ($\lambda$5697) | N IV ($\lambda 4058$, $\lambda\lambda 7109,7122$)
($3.9\times 10^{8}-3.1\times 10^{9}$ $\mathrm{L_{\odot}}$) | C IV ($\lambda\lambda 5801,5811$) |
| N III ($\lambda\lambda 4634,4640$) | N III ($\lambda\lambda 4634,4640$)
High luminosity ($>6.3\times 10^{9}$ $\mathrm{L_{\odot}}$) | O VI ($\lambda\lambda 3811,3834$) | Lack of O VI ($\lambda\lambda 3811,3834$)
Figure 4: The day 1.36 spectrum (black) compared with three different surface
abundance models from Boian & Groh (2019) with $L=1.5\times 10^{9}$
$\mathrm{L_{\odot}}$ and $\dot{M}=3\times 10^{-3}$ M⊙$~{}\mathrm{yr^{-1}}$.
The solar abundance represents a low-mass RSG (yellow), CNO-processed
abundance represents a high-mass RSG/BSG/YSG (blue), and He-rich abundance
represents a LBV/WN/stripped star (pink). Vertical dashed lines represent the
ions that are identified in the observed spectra that are labeled at the top
of the plot. Vertical dotted lines in light gray show ions that are not
detected but are part of the progenitor diagnostics detailed in Boian & Groh
(2019). We convolve the model spectra with a Gaussian kernel to mimic the
resolution of the observed spectra. While some individual features in each
model are matched by the observed spectrum, there are differences between the
model and observed spectrum that indicate that the conditions of the model are
not the same as the observations.

###  5.2 Dessart & Hillier models

In another study, Dessart et al. (2017) model the spectroscopic signature of
CSM interaction with a variety of RSG mass-loss rates
($\dot{M}=10^{-6}-10^{-2}$ M⊙ $\mathrm{yr^{-1}}$) and atmospheric density
scale heights ($H_{\rho}=0.01,0.1,0.3R_{*}$). The base of each RSG model is a
15 M⊙ star onto which they add an atmosphere with a given density scale
height, which transitions to wind mass loss when the density of the atmosphere
equals the density of the wind. The wind is then extended to
$R_{\mathrm{out,CSM}}=5\times 10^{14}$ cm for all but one model which is
extended to $R_{\mathrm{out,CSM}}=2\times 10^{14}$ cm. At this point all
models transition to an $\dot{M}=10^{-6}$ M⊙ $\mathrm{yr^{-1}}$ wind. The
parameters of each model are summarized in Table 2. Each of these models is
evolved from shock breakout to over 10 days and snapshots of the spectra and
light curves are reported.

Given the uncertainties in the explosion epoch for the observations and
challenges of explosion in the models (e.g. core-collapse vs shock breakout,
varying shock breakout time scales in dense CSM), we compare our observed
spectra with the full suite of models for all mass-loss rates and epochs. In
the weak-wind models (r1w1, r1w1h, r2w1), the only spectrum in the time series
with narrow emission features is the first spectrum at shock break out. As our
lines are clearly present throughout the first 5–7 days of evolution, we do
not examine these models further. With this cut alone, we constrain the mass-
loss rate to be $\dot{M}>10^{-3}$$\mathrm{M_{\odot}\,yr^{-1}}$. The remaining
models (r1w4, r1w6, r1w5r, and r1w5h) show multiple epochs of narrow emission
features. These features give way to narrow P Cygni profiles in the He II
($\lambda 4685.5$) and $\mathrm{H\alpha}$ lines which are eventually replaced
by broader features originating from the bulk motion of the ejecta.

Again, none of the models reproduces the observed spectra. Broadly, the
emission features in the model spectra fade much more rapidly than in the
observed spectra, if they are present at all. This implies that the CSM in SN
2023ixf extends beyond that of the models or has a higher density (e.g. the
r1w6 model is the only one with narrow emission lines that persist past 2
days). Particularly challenging is the blended N III ($\lambda\lambda
4634.0,4640.6$)/C III ($\lambda\lambda 4647.5,4650.0$) complex and the He I
lines, which are clearly visible in the observed spectrum at day 1.36 and
quickly fades below detection by day 2.0. These lines are not present in most
of the model spectra at any epoch. Additionally, the spectra of SN 2023ixf
never show N V ($\lambda 4610$), which is present in the early spectra of all
of the strong wind models, although in some models this has faded by the phase
of our first spectrum (e.g. r1w4 in Figure 5). Instead, we see N III ($\lambda
4636$), which has a similar flux to He II ($\lambda 4686$) in the
classification spectrum and then fades over the subsequent 0.5d. Although the
r1w4 and r1w6 models show this feature just prior to the development of the
narrow P Cygni features, it is always significantly weaker than the He II
emission. Given the rapid evolution of this feature, this could be a function
of the model sampling. We note that we do not see evidence of the rise of N V
described by Jacobson-Galan et al. (2023) in our spectra. Rather, we find the
asymmetric blue wings of He II ($\lambda 4685.5$) to be more consistent with N
III ($\lambda\lambda 4634.0,4640.6$)/C III ($\lambda\lambda 4647.5,4650.0$).

We find the best agreement with the r1w4 and r1w6 models which, while not able
to reproduce the observed line ratios, show N III, He I ($\lambda 7065$), C IV
($\lambda\lambda 5801,5811$), and N IV ($\lambda 4057$, $\lambda 7122$)
features. Although present, the C IV and N IV are significantly stronger in
the first model spectra than in the observed spectra. Additionally, the
strength of the Balmer emission lines is better matched in these models as is
the lack of the O V ($\lambda 5597$). Figure 5 shows the spectral evolution of
SN 2023ixf compared to the r1w4 model. While the first spectrum matches well,
the narrow lines disappear from the model by day 2 and narrow P Cygni profiles
emerge. In the final spectrum, the model shows significantly more
$\mathrm{H\alpha}$ emission than the observation. Interestingly, the mass-loss
rate of this model is $10^{-3}$ M⊙ $\mathrm{yr^{-1}}$, consistent with the
conclusions from the comparison to the models of Boian & Groh (2019).

The first r1w6 model spectrum on day 1.3 clearly shows N V. Additionally, r1w6
model time series does not show the N III shoulder on He II until 1-2 days
after it has disappeared from the observed spectra, however, the features are
otherwise reasonably matched. The disappearance of the narrow emission lines
and evolution of $\mathrm{H\alpha}$ are better matched in the r1w6 model. This
is consistent with the findings of Jacobson-Galan et al. (2023), who find a
best fit model in their custom grid is the r1w6 model with a larger radius.
Given the challenges of reproducing the observed spectra with published
models, we conservatively conclude that the mass-loss rate for the progenitor
of SN 2023ixf was between 10-3 and 10-2 $\mathrm{M_{\odot}\,yr^{-1}}$, based
on the presence of persistent lines after the initial shock breakout spectrum.
The spectra themselves are sensitive to density, which is parameterized as
$\dot{M}=\rho v_{w}$, thus a different mass-loss rate is inferred if the
assumed velocity is different. The wind velocity in these models is
$v_{w}=50~{}\mathrm{km\,s^{-1}}$, which is about a factor of 3 smaller than
the measured velocity. However, given the order of magnitude range in mass-
loss rate inferred from these models, we do not further modify the mass-loss
rates.

Table 2: CSM Parameters for the Models of Dessart et al. (2017)

Model Name | Mass-loss Rate ($\mathrm{M_{\odot}\,yr^{-1}}$) | Radius ($\mathrm{R_{\odot}}$) | Scale Height ($R_{\star}$) | Transition Radius⋆ (cm)
---|---|---|---|---
r1w1 | $10^{-6}$ | 501 | 0.01 | $5\times 10^{14}$
r1w2 | $10^{-5}$ | 501 | 0.01 | $5\times 10^{14}$
r1w3 | $10^{-4}$ | 501 | 0.01 | $5\times 10^{14}$
r1w4 | $10^{-3}$ | 501 | 0.01 | $5\times 10^{14}$
r1w5 | $5\times 10^{-3}$ | 501 | 0.01 | $5\times 10^{14}$
r1w6 | $10^{-2}$ | 501 | 0.01 | $5\times 10^{14}$
r1w1h | power law† | 501 | 0.3 | $5\times 10^{14}$
r1w5r | $10^{-5}$ | 501 | 0.01 | $2\times 10^{14}$
r2w1 | $10^{-6}$ | 1107 | 0.01 | $5\times 10^{14}$

Note. — †power-law exponent with an exponent of 12

Note. — ⋆Radius at which the density transitions to $\dot{M}=10^{-6}$
$\mathrm{M_{\odot}\,yr^{-1}}$

Figure 5: A comparison of the spectral evolution of SN 2023ixf (black) to the
r1w4 model of Dessart et al. (2017) (pink). Phase are shown to the right of
the figure in pink for the model and black from the observed spectra. While
the first observation at 1.2 day matches fairly well, the model spectra evolve
much more rapidly, with P Cygni profiles developing at day 2 and all emission
disappearing by day 4, while the observations show emission through day 6-7.
The model spectrum at 0.5d shows a number of lines that have faded by 1d (e.g.
N V).

##  6 Model-independent CSM Constraints

The properties and evolution of the narrow emission lines allow us to compute
order of magnitude estimates of the CSM characteristics. A lower limit on the
outer CSM radius can be calculated from the brightest $\mathrm{H\alpha}$ flux
($F_{\mathrm{H\alpha}}$; Yaron et al., 2017; Ofek et al., 2013). Briefly, the
$\mathrm{H\alpha}$ luminosity ($L_{\mathrm{H\alpha}}$) is produced by the
recombination of ionized H. Assuming a spherically symmetric CSM that is
composed of hydrogen, all of which is ionized, we can use the
$\mathrm{H\alpha}$ luminosity to calculate the total CSM mass. The total mass
can also be calculated by integrating a constant velocity wind density profile
over the radial extent of the CSM. Equating these two relations allows us to
determine the radius of the CSM:

$r\gtrsim\frac{\kappa L_{\mathrm{H\alpha}}}{A}$ (1)

where $\kappa=0.34~{}\mathrm{cm^{2}\,g^{-1}}$ is the electron scattering
opacity of the CSM. We use the distance $D$ to convert $F_{\mathrm{H\alpha}}$
to $L_{\mathrm{H\alpha}}$ via $L_{\mathrm{H\alpha}}=4\pi
D^{2}F_{\mathrm{H\alpha}}$. A is defined as

$A=\frac{4\pi h\nu_{H}\alpha_{H}^{\text{eff}}}{\mu_{p}m_{p}^{2}}$ (2)

where $\nu_{H}=4.56\times 10^{14}~{}\mathrm{Hz}$ is the frequency of
$\mathrm{H\alpha}$, $\alpha_{H}^{\text{eff}}=8.7\times
10^{14}~{}\mathrm{cm^{3}\,s^{-1}}$ is the H recombination coefficient for case
B recombination at $T_{\text{eff}}=10000~{}\mathrm{K}$, $\mu_{p}=0.5$ is the
mean molecular weight, $m_{p}$ is the proton mass, and $h$ is Planck’s
constant. We treat Equation 1 as a lower limit on the outer CSM radius, as
either the composition or ionization assumption may not be true, which would
lead to a larger CSM radius than calculated here. Additionally, we assume that
the CSM above the emitting region is transparent to the $\mathrm{H\alpha}$
photons. If this is not true, it would also lead to a larger CSM radius.

To measure the flux of $\mathrm{H\alpha}$, we fit a blackbody to the continuum
and subtract it from the flux of our 1.36 day spectrum, which has the maximum
$\mathrm{H\alpha}$ flux of our spectral series. To the continuum-subtracted
flux, we simultaneously fit broad and narrow Lorentzian emission profiles. We
integrate this fit from 6300–6800 Å to find $F_{\mathrm{H\alpha}}=3.18\times
10^{-13}~{}\mathrm{erg\,cm^{-2}\,s^{-1}}$. Using a distance of 6.85 Mpc, we
find $L_{\mathrm{H\alpha}}=1.78\times 10^{39}~{}\mathrm{erg\,s^{-1}}$ and
$R_{\text{CSM, out}}\gtrsim 8.7\times 10^{13}~{}\mathrm{cm}$. Assuming a
spherical wind with a constant velocity, mass-loss rate, and homogeneous
density structure, the density can be calculated from the radius:

$\rho=\frac{1}{\kappa\,r}$ (3)

From this, we calculate a density of $3.4\times
10^{-14}~{}\mathrm{g\,cm^{-3}}$. Assuming a typical RSG wind of $v_{w}\approx
10~{}\mathrm{km\,s^{-1}}$, this mass loss event would have begun $\gtrsim
3~{}\mathrm{yr}$ before explosion. However, using the mass-loss rate derived
from high-resolution spectroscopy in Smith et al. (2023) of $v_{w}\approx
150~{}\mathrm{km\,s^{-1}}$ we find a much smaller start time of $\sim 2$
months prior to explosion.

The narrow features in the CSM are only present when there is unshocked,
photoionized CSM in front of the SN shock and ejecta. Therefore, we expect
these features to disappear when the material producing them is swept up by
the ejecta, and we can use this information to calculate the radius of the
CSM. In SN 2023ixf, the unshocked, narrow features disappear from
$\mathrm{H\alpha}$ 3–4 days after explosion. Smith et al. (2023) find this
corresponds to a radius of $R=(3-5)\times 10^{14}$ cm. However, at this phase,
there is still material in front of the photosphere which produces
intermediate width lines and eventually an intermediate width P Cygni profile
in $\mathrm{H\alpha}$. These intermediate width lines disappear around day
$\sim$6–7. On day 13, clear broad absorption is visible in H$\beta$ and from
this we approximate an average ejecta velocity of $\sim
9000~{}\mathrm{km\,s^{-1}}$. Putting this together with the time that the
lines disappear, we calculate a CSM radius of $R_{\mathrm{CSM,out}}\sim
5.4\times 10^{14}~{}\mathrm{cm}$. Interestingly, this is exactly where the
models of Dessart et al. (2017) transition to $\mathrm{\dot{M}=10^{-6}}$
$\mathrm{M_{\odot}\,yr^{-1}}$. Again, assuming a constant RSG wind of
$v_{w}\approx 10~{}\mathrm{km\,s^{-1}}$ (or $v_{w}\approx
150~{}\mathrm{km\,s^{-1}}$), this implies that the event began $\sim
17\,\mathrm{yr}$ ($1\,\mathrm{yr}$ for $v_{w}\approx
150~{}\mathrm{km\,s^{-1}}$) before explosion. Assuming a constant wind density
profile, the density is related to radius by:

$\rho=\frac{\dot{M}}{4\pi v_{w}R_{\mathrm{CSM}}^{2}}$ (4)

Using a representative mass-loss rate of $5\times 10^{-3}$
$\mathrm{M_{\odot}\,yr^{-1}}$, we find a density of $\rho=8.5\times
10^{-14}~{}\mathrm{g\,cm^{-3}}$ ($\rho=5.6\times
10^{-14}~{}\mathrm{g\,cm^{-3}}$ for $v_{w}\approx 150~{}\mathrm{km\,s^{-1}}$).
This is consistent with the lower limit calculated from the $\mathrm{H\alpha}$
luminosity.

We compare the mass-loss rates derived in Section 5 to a population of Type II
SNe with flash ionization features to identify how unusual (or normal) these
values are. We use the sample of Boian & Groh (2020), who use CMFGEN to model
a sample of 17 Type II SNe that show narrow emission features in their early
spectra, indicating CSM interaction. In Section 5, we used a spectrum taken
1.36 days after explosion to determine the mass-loss rate and SN luminosity
for SN 2023ixf. Given our extensive time sampling, we investigate whether the
epoch of the spectrum used to do this would affect which model was selected.
We tested this using spectra taken on days 2.26, 3.16, and 5.15. Given the
uncertainty in the epoch of the models, we consider all model epochs for each
observed spectrum. We find that, in all cases, regardless of the epoch used,
we would select the models and thus mass-loss rates from Section 5.
Additionally, we find that with the day 5.15 spectrum, we would include the
r1w5h model, which has a mass-loss rate between the two models we selected
(r1w4 and r1w6). This demonstrates that the differences between the models are
robust to observed epoch as long as narrow features are present.

With this confirmation, we add SN 2023ixf to Figure 8 from Boian & Groh
(2020), which shows the SN luminosity and mass-loss rate for all SNe in their
sample (Figure 6). We find the mass-loss rate of SN 2023ixf to be fairly low
when compared to the range of mass-loss rates for other SN with early
interaction, although it is in no way an outlier. This is consistent with a
lower density CSM causing the persistence of the higher-ion levels of N IV and
C IV.

While in line with other interacting events, the mass-loss rates determined
for the confined CSM in this paper are significantly higher than steady-state
RSG winds. Using the progenitor luminosity
$\log(L/$$\mathrm{L_{\odot}}$$\mathrm{)=4.94}$ found by Jencson et al. (2023)
and the de Jager et al. (1988) relation, the expected mass-loss rate would be
$\dot{M}\approx 10^{-5.7}$ $\mathrm{M_{\odot}\,yr^{-1}}$ and even lower for
the progenitor luminosity identified in Kilpatrick et al. (2023). Using the
relationship of Beasor et al. (2020) and the stellar parameters of Jencson et
al. (2023), the expected mass-loss rate would be $\dot{M}\approx 2\times
10^{-6}$ $\mathrm{M_{\odot}\,yr^{-1}}$. Even using the enhanced mass-loss
rates of Ekström et al. (2012), the mass-loss rate would only be
$\dot{M}\approx 10^{-5.2}$ $\mathrm{M_{\odot}\,yr^{-1}}$. This implies that
the mass-loss event which led to this CSM was vastly greater than the nominal
RSG mass loss. On the other hand, the largest mass-loss rate that we find is
consistent with the lower range of mass-loss rates required by Morozova et al.
(2018) (assuming a 10 $\mathrm{km\,s^{-1}}$ wind) to fit the rapid light curve
rises seen in Type II SNe.

Figure 6: The SN luminosity and mass-loss rate derived by Boian & Groh (2020)
for a sample of 17 SNe compared to the best model of SN 2023ixf (black). Upper
limits on mass loss are shown with semi-transparent markers and arrows, while
determined values are solid. We determine a scaled mass-loss rate of
$\dot{M}\approx 4.5\times 10^{-3}$ $\mathrm{M_{\odot}\,yr^{-1}}$ and a scaled
luminosity of $L\approx 2.6\times 10^{9}$$\mathrm{L_{\odot}}$ from the models
of Boian & Groh (2019) (with slight variation due to different surface
abundances) which we plot as a black stars. In practice, these are located at
virtually the same location on this plot. We also include the mass-loss rates
of the r1w4 and r1w6 models of Dessart et al. (2017) ($\dot{M}=10^{-3}$
$\mathrm{M_{\odot}\,yr^{-1}}$ and $\dot{M}=10^{-2}$
$\mathrm{M_{\odot}\,yr^{-1}}$, respectively) which are shown as black
triangles. Dessart et al. (2017) does not explore variations in luminosity and
we therefore use the luminosity derived from Boian & Groh (2019) solar
abundance models and do not scale the mass-loss rates. The mass-loss rate of
SN 2023ixf is in line with the lower end of the range of mass-loss rates in
this sample.

We make a final note on the first 0.5d evolution of the N III ($\lambda\lambda
4634.0,4640.6$)/C III ($\lambda\lambda 4647.5,4650.0$) emission line. If the
CSM were suddenly ionized by a high energy flash from shock breakout photons,
then one would expect that the CSM is almost instantaneously (modulo the light
travel time) ionized. Thus, as time passes in the days after the initial flash
ionization, the spectral evolution should proceed from high ionization to low
ionization species as the initially highly ionized CSM recombines over a
timescale determined by the density of the CSM. However, we observe the
opposite. Spectra of SN 2023ixf show that relatively low-ionization levels are
present in the earliest epochs (i.e. narrow N III ($\lambda\lambda
4634.0,4640.6$)/C III ($\lambda\lambda 4647.5,4650.0$) and He I emission on
day 1-2). Over the following day, these fade away as other higher ionization
features (such as He II ($\lambda 4685.5$), N iv, C iv) strengthen. This is
also discussed by Smith et al. (2023) and Jacobson-Galan et al. (2023). This
implies that we are witnessing a gradual or delayed ionization of the CSM, an
effect never observed before. This effect seems inconsistent with a sudden
flash ionization from shock breakout alone, and more indicative of a slowly
varying source of ionization (i.e. radiation from the CSM interaction shock
itself or asymmetry in the ejecta).

##  7 Summary & Conclusions

The comprehensive spectroscopic data set presented in this paper provides sub-
day cadence spectroscopic observations for the first week of evolution and at
least daily cadence through day 14. To our knowledge, no other SN with flash
ionization features has been studied with this cadence for this length of
time. The unique combination of cadence and duration of these observations
will enable the community to trace the CSM density profile, and therefore the
mass-loss history of the progenitor, in a way that has never before been
possible.

These spectra show narrow high-ionization emission features with Lorentzian
wings. We identify emission from H I, He I, He II, C III, C IV, O III, N III,
and N IV. He I, C III, N III, and O III fade over our first 0.5d of monitoring
($\sim$1.5-2d since explosion). We find the spectrum at 1.36d most closely
resembles that of SNe 2017ahn and 2020pni and looks significantly different
from SN 2013fs. Over time, the evolution closely matches that of SN 2014G, as
the high-ionization features fade and eventually so does He II and H I. The
differences in the evolution of these different SNe imply that SN 2023ixf has
a lower density CSM than SNe 2017ahn and 2020pni and a smaller radial extent.
By day 13, broad P Cygni profiles have developed indicating emission from the
SN ejecta.

We compare the same 1.36d spectrum with the models of Boian & Groh (2019),
finding the spectrum most closely resembles the models with
$L\approx$2.6$\times 10^{9}$ $\mathrm{L_{\odot}}$ and
$\dot{M}\approx$4.5$\times 10^{-3}$ $\mathrm{M_{\odot}\,yr^{-1}}$, although we
do not find any model that reproduces the line ratios in our observed spectra
and therefore cannot use these diagnostics to infer a progenitor mass.

We also relate the full spectral evolution over the first two weeks to the
models of Dessart et al. (2017), which examine different mass-loss rates and
atmospheric scale heights. We find the spectra of SN 2023ixf are best
represented by the r1w4 and r1w6 models, corresponding to mass-loss rates of
$\mathrm{\dot{M}=10^{-3}}$ M⊙ $\mathrm{yr^{-1}}$ and
$\mathrm{\dot{M}=10^{-2}}$ M⊙ $\mathrm{yr^{-1}}$ respectively. However, we
note that in r1w4 model, the narrow emission lines disappear much more rapidly
from the model than we observe indicating that a larger radial extent is
required in the model. We find that despite the rapid evolution of the
spectrum over the first five days, the spectra are most consistent with the
same models, regardless of epoch used to identify them.

Finally, we use the narrow lines to calculate the properties of the CSM. Using
the maximum $\mathrm{H\alpha}$ flux and assuming a spherical geometry, we find
$R_{\text{CSM, out}}\gtrsim 8.7\times 10^{13}~{}\mathrm{cm}$, implying the CSM
ejection began at least 3 years ago if it is expanding at 10
$\mathrm{km\,s^{-1}}$ and 2 months ago if the observed velocity of
$150~{}\mathrm{km\,s^{-1}}$. At this radius, we find a CSM density of
$3.3\times 10^{-14}~{}\mathrm{g\,cm^{-3}}$. Using the epoch at which time the
narrow emission features disappear, we find a consistent radius of
$R_{\text{CSM}}=5.4\times 10^{14}~{}\mathrm{cm}$. With this radius and RSG
wind $v_{w}\approx 10~{}\mathrm{km\,s^{-1}}$ ($v_{w}\approx
150~{}\mathrm{km\,s^{-1}}$), the CSM ejection began 17 years ago (1 year ago)
and the density is $\rho=8.5\times 10^{-14}~{}\mathrm{g\,cm^{-3}}$
($\rho=5.6\times 10^{-14}~{}\mathrm{g\,cm^{-3}}$). Comparing SN 2023ixf to a
sample of 17 SNe with early CSM interaction, we find the mass-loss rate in
line with the lower end of the distribution and the low luminosity. We note
that this analysis assumes spherically symmetric CSM. Asymmetric CSM (such as
that proposed by Smith et al. (2023) or clumped CSM (e.g. Dessart et al.,
2018) would alter these conclusions, although the details of the effect will
depend on the exact configuration.

SN 2023ixf is an extraordinary SN, combining proximity with early detection
and classification, and tight constraints on explosion. The immediate
announcement of the discovery and classification allowed us to harness our
resources and observe the detailed evolution of the CSM interaction over the
first two weeks. With these observations we identify a significantly higher
mass-loss rate than the nominal RSG mass-loss rate of
$\mathrm{\dot{M}=10^{-6}}$ M⊙ $\mathrm{yr^{-1}}$. This indicates either a
superwind or period of eruptive mass loss (but see also Kochanek, 2019). While
both the models of Boian & Groh (2019) and Dessart et al. (2017) are unable to
match the temporal evolution and the relative flux ratios, the majority of the
emission lines present are reproduced and we find the prospects of a custom
model to match the observations encouraging. The spectroscopic data set
presented in this paper can help guide future modeling efforts and be used to
benchmark the evolution flash features in any SN.

## Acknowledgments

We thank our anonymous referee for providing comments to improve the paper.
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 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. Time domain research by the University of Arizona
team and D.J.S. is supported by NSF grants AST-1821987, 1813466, 1908972, &
2108032, and by the Heising-Simons Foundation under grant #20201864. The
research by Y.D., S.V., N.M., and E.H. is supported by NSF grants AST-2008108.
J.E.A. is supported by the international Gemini Observatory, a program of
NSF’s NOIRLab, which is managed by the Association of Universities for
Research in Astronomy (AURA) under a cooperative agreement with the National
Science Foundation, on behalf of the Gemini partnership of Argentina, Brazil,
Canada, Chile, the Republic of Korea, and the United States of America. The
research of J.C.W. and J.V. is supported by NSF AST-1813825. J.V. is also
supported by OTKA grant K-142534 of the National Research, Development and
Innovation Office, Hungary. L.S. and M.W.C. acknowledge support from the
National Science Foundation with grant numbers PHY-2010970 and OAC-2117997.
A.Z.B acknowledges support from the European Research Council (ERC) under the
European Union’s Horizon 2020 research and innovation programme (Grant
agreement No. 772086). AR acknowledges support from ANID BECAS/DOCTORADO
NACIONAL 21202412. We thank the MMT director, G. Williams, for granting
Director’s Discretionary Time for the Hectospec spectral sequence. This paper
made use of the modsCCDRed data reduction code developed in part with funds
provided by NSF Grants AST-9987045 and AST-1108693. A.P. and P.O. acknowledge
support of the PRIN-INAF 2022 project “Shedding light on the nature of gap
transients: from the observations to the models”. The SNICE research group
acknowledges financial support from the Spanish Ministerio de Ciencia e
Innovación (MCIN), the Agencia Estatal de Investigación (AEI)
10.13039/501100011033, the European Social Fund (ESF) “Investing in your
future”, the European Union Next Generation EU/PRTR funds, the Horizon 2020
Research and Innovation Programme of the European Union, and by the Secretary
of Universities and Research (Government of Catalonia), under the
PID2020-115253GA-I00 HOSTFLOWS project, the 2019 Ramón y Cajal program
RYC2019-027683-I, the 2021 Juan de la Cierva program FJC2021-047124-I, the
Marie Skłodowska-Curie and the Beatriu de Pinós 2021 BP 00168 programme, and
from Centro Superior de Investigaciones Científicas (CSIC) under the PIE
project 20215AT016, and the program Unidad de Excelencia María de Maeztu
CEX2020-001058-M.

We thank David Bohlender, Dmitry Monin, and James Di Francesco for obtaining
DAO spectra and the whole LBT team, especially Alexander Becker and Jennifer
Power. Based on observations obtained with the Hobby-Eberly Telescope (HET),
which is a joint project of the University of Texas at Austin, the
Pennsylvania State University, Ludwig-Maximillians-Universitaet Muenchen, and
Georg-August-Universitaet Goettingen. The HET is named in honor of its
principal benefactors, William P. Hobby and Robert E. Eberly. The Low
Resolution Spectrograph 2 (LRS2) was developed and funded by the University of
Texas at Austin McDonald Observatory and Department of Astronomy, and by
Pennsylvania State University. We thank the Leibniz-Institut fur Astrophysik
Potsdam (AIP) and the Institut fur Astrophysik Goettingen (IAG) for their
contributions to the construction of the integral field units. The Liverpool
Telescope is operated on the island of La Palma by Liverpool John Moores
University in the Spanish Observatorio del Roque de los Muchachos of the
Instituto de Astrofisica de Canarias with financial support from the UK
Science and Technology Facilities Council. The data presented here were
obtained in part with ALFOSC, which is provided by the Instituto de
Astrofisica de Andalucia (IAA) under a joint agreement with the University of
Copenhagen and NOT. Based on observations made with the Nordic Optical
Telescope, owned in collaboration by the University of Turku and Aarhus
University, and operated jointly by Aarhus University, the University of Turku
and the University of Oslo, representing Denmark, Finland and Norway, the
University of Iceland and Stockholm University at the Observatorio del Roque
de los Muchachos, La Palma, Spain, of the Instituto de Astrofisica de
Canarias. This project has received funding from the European Union’s Horizon
2020 research and innovation programme under grant agreement No 101004719
(ORP: OPTICON RadioNet Pilot). Observations reported here were obtained at the
MMT Observatory, a joint facility of the University of Arizona and the
Smithsonian Institution. Based on observations made with the 1.22-m Galileo
Galilei Telescope of the Padova University in the Asiago site. This paper made
use of the modsCCDRed data reduction code developed in part with funds
provided by NSF Grants AST-9987045 and AST-1108693.

##  Appendix A Spectroscopic Observations

Table 3 lists the date, telescope, instrument, and resolving power for each
spectroscopic observation used in this paper.

Table 3: Log of Spectroscopic Observations

Phase (d) | Time | MJD | Telescope | Instrument | R ($\lambda$/$\Delta\lambda$)
---|---|---|---|---|---
1.18 | 2023-05-19 22:23:45 | 60083.93 | LT | SPRAT | 350
1.36 | 2023-05-20 02:39:56 | 60084.11 | NOT | ALFOSC | 360
1.54 | 2023-05-20 07:03:24 | 60084.29 | Bok | B&C | 700
1.67 | 2023-05-20 10:00:43 | 60084.42 | MMT | Hectospec | 1325
2.13 | 2023-05-20 21:04:19 | 60084.88 | HCT | HFOSC | 350
2.26 | 2023-05-21 00:14:57 | 60085.01 | NOT | ALFOSC | 360
2.43 | 2023-05-21 04:19:39 | 60085.18 | MMT | Hectospec | 1325
2.50 | 2023-05-21 06:07:11 | 60085.25 | FTN | FLOYDS | 500
2.53 | 2023-05-21 06:40:49 | 60085.28 | Bok | B&C | 700
2.66 | 2023-05-21 09:50:44 | 60085.41 | MMT | Hectospec | 1325
2.76 | 2023-05-21 12:07:58 | 60085.51 | FTN | FLOYDS | 500
3.09 | 2023-05-21 20:03:50 | 60085.84 | Galileo | B&C | 2762
3.15 | 2023-05-21 21:41:41 | 60085.90 | LT | SPRAT | 350
3.16 | 2023-05-21 21:48:37 | 60085.91 | NOT | ALFOSC | 360
3.39 | 2023-05-22 03:21:44 | 60086.14 | LBT | MODS | 2075
3.42 | 2023-05-22 03:57:46 | 60086.17 | Bok | B&C | 700
3.54 | 2023-05-22 07:00:07 | 60086.29 | FTN | FLOYDS | 500
3.61 | 2023-05-22 08:41:45 | 60086.36 | Bok | B&C | 700
3.64 | 2023-05-22 09:15:15 | 60086.39 | MMT | Hectospec | 1325
4.02 | 2023-05-22 18:31:41 | 60086.77 | HCT | HFOSC | 350
4.19 | 2023-05-22 22:33:56 | 60086.94 | NOT | ALFOSC | 300
4.19 | 2023-05-22 22:32:09 | 60086.94 | Galileo | B&C | 1195
4.24 | 2023-05-22 23:50:00 | 60086.99 | Other† | Other† | -†
4.41 | 2023-05-23 03:43:23 | 60087.16 | MMT | Hectospec | 1325
4.45 | 2023-05-23 04:47:09 | 60087.20 | Bok | B&C | 700
5.10 | 2023-05-23 20:16:48 | 60087.85 | Galileo | B&C | 983
5.15 | 2023-05-23 21:33:06 | 60087.90 | NOT | ALFOSC | 300
5.20 | 2023-05-23 22:48:32 | 60087.95 | LT | SPRAT | 350
5.25 | 2023-05-24 00:01:26 | 60088.00 | Galileo | B&C | 983
5.41 | 2023-05-24 03:47:15 | 60088.16 | MMT | Hectospec | 1325
5.51 | 2023-05-24 06:07:28 | 60088.26 | Bok | B&C | 700
6.35 | 2023-05-25 02:18:49 | 60089.10 | NOT | ALFOSC | 300
6.43 | 2023-05-25 04:13:21 | 60089.18 | MMT | Hectospec | 1325
6.52 | 2023-05-25 06:35:59 | 60089.27 | Bok | B&C | 700
6.58 | 2023-05-25 08:02:16 | 60089.33 | FTN | FLOYDS | 500
7.43 | 2023-05-26 04:15:36 | 60090.18 | MMT | Hectospec | 1325
7.50 | 2023-05-26 05:55:24 | 60090.25 | FTN | FLOYDS | 500
7.60 | 2023-05-26 08:17:43 | 60090.35 | Bok | B&C | 700
8.43 | 2023-05-27 04:21:08 | 60091.18 | Bok | B&C | 700
8.73 | 2023-05-27 11:25:38 | 60091.48 | FTN | FLOYDS | 500
9.11 | 2023-05-27 20:35:31 | 60091.86 | Galileo | B&C | 2455
9.45 | 2023-05-28 04:52:00 | 60092.20 | Bok | B&C | 700
9.53 | 2023-05-28 06:40:09 | 60092.28 | FTN | FLOYDS | 500
9.61 | 2023-05-28 08:35:03 | 60092.36 | APO | ARCES | 30000
10.36 | 2023-05-29 02:47:20 | 60093.12 | APO | ARCES | 30000
10.73 | 2023-05-29 11:26:52 | 60093.48 | FTN | FLOYDS | 500
11.30 | 2023-05-30 01:11:14 | 60094.05 | LT | SPRAT | 350
11.48 | 2023-05-30 05:24:06 | 60094.23 | Bok | B&C | 700
11.75 | 2023-05-30 12:00:24 | 60094.50 | FTN | FLOYDS | 500
12.10 | 2023-05-30 20:18:14 | 60094.85 | Galileo | B&C | 2007
12.18 | 2023-05-30 22:18:38 | 60094.92 | INT | IDS | 1092
12.55 | 2023-05-31 07:06:58 | 60095.30 | FTN | FLOYDS | 500
13.15 | 2023-05-31 21:40:19 | 60095.90 | Galileo | B&C | 1019
13.58 | 2023-06-01 07:53:58 | 60096.33 | FTN | FLOYDS | 500
14.16 | 2023-06-01 21:51:50 | 60096.91 | Galileo | B&C | 1037
14.18 | 2023-06-01 22:19:25 | 60096.93 | LT | SPRAT | 350
14.53 | 2023-06-02 06:44:48 | 60097.28 | HET | LRS2 | 600
15.26 | 2023-06-03 00:19:09 | 60098.01 | INT | IDS | 1092

Note. — † This spectrum on TNS does not have any information regarding
telescope, instrument, or resolving power.

##  Appendix B Line Identification

Table 4 gives the ion and wavelength of the lines identified and shown
throughout this paper.

Table 4: List of ions observed in the day 1.36 spectrum.  Ion | Wavelength (Å)
---|---
He I | 3888.64
O III | 3961.59
H I | 3970.0
C III | 4056.0
N IV | 4057.76
N III | 4097.33
H I | 4101.73
He II | 4200.0
H I | 4340.47
O III | 4519.62
O III | 4540.4
N III | 4634.0
N III | 4640.64
C III | 4647.5
C III | 4650.0
He II | 4685.5
N III | 4858.82
He II | 4860.0
H I | 4861.3
O III | 5006.8
He II | 5412.0
O III | 5592.37
C III | 5695.9
C IV | 5801.3
C IV | 5811.98
He I | 5875.6
He II | 6559.8
H I | 6562.7
He I | 6678.15
C III | 7037.25
He I | 7065.19
N IV | 7103.24
N IV | 7109.0
N IV | 7122.98
He I | 7281.34

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