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# SN 2023ixf in the Pinwheel Galaxy M101: From Shock Breakout to the Nebular Phase

WeiKang Zheng¹, Luc Dessart², Alexei V. Filippenko¹, Yi Yang (杨轶)¹,³, Thomas G. Brink¹, Thomas de Jaeger⁴, Sergiy S. Vasylyev¹, Schuyler D. Van Dyk⁵, Kishore C. Patra¹,⁶, Wynn V. Jacobson-Galán¹,⁷, Gabrielle E. Stewart¹, Efrain Alvarado, III¹,⁸, Veda Arikatla¹, Pallas Beddow¹, Andreas Betz¹, Emma Born¹, Kate Bostow¹,⁸, Adam J. Burgasser⁹, Osmin Caceres¹⁰, Evan M. Carrasco⁶, Elma Chuang¹, Asia DeGraw¹, Elinor L. Gates¹¹, Eli Gendreau-Distler¹, Cooper Jacobus¹, Connor Jennings¹,¹², Preethi R. Karpoor⁹, Paul Lynam¹¹, Ann Mina¹, Katherine Mora¹, Neil Pichay¹, Jyotsna Ravi¹, Jon Rees¹¹, R. Michael Rich¹⁰, Sophia Risin¹, Nathan R. Sandford¹³, Alessandro Savino¹, Emma Softich⁹, Christopher A. Theissen⁹, Edgar P. Vidal¹,¹⁴, William Wu¹, and Yoomee Zeng¹

¹Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA; weikang@berkeley.edu, afilippenko@berkeley.edu, yi_ yang@mail.tsinghua.edu.cn

²Institut d'Astrophysique de Paris, CNRS-Sorbonne Université, 98 bis boulevard Arago, F-75014 Paris, France; dessart@iap.fr

³Physics Department, Tsinghua University, Beijing, 100084, People's Republic of China

⁴LPNHE (CNRS/IN2P3, Sorbonne Université, Université Paris Cité), Laboratoire de Physique Nucléaire et de Hautes Énergies, 75005, Paris, France; dejaeger.thomas@gmail.com

⁵Caltech/IPAC, Mail Code 100-22, Pasadena, CA 91125, USA

⁶Department of Astronomy & Astrophysics, University of California, Santa Cruz, CA 95064, USA

⁷Department of Astronomy and Astrophysics, California Institute of Technology, Pasadena, CA 91125, USA

⁸Department of Physics and Astronomy, San Francisco State University, 1600 Holloway Avenue, San Francisco, CA 94132, USA

⁹Department of Astronomy & Astrophysics, University of California at San Diego, La Jolla, CA 92093, USA

¹⁰Department of Physics and Astronomy, University of California, Los Angeles, CA 90095-1547, USA

¹¹University of California Observatories/Lick Observatory, Mount Hamilton, CA 95140, USA

¹²Department of Astronomy, Yale University, New Haven, CT 06511, USA

¹³Department of Astronomy and Astrophysics, University of Toronto, 50 St. George Street, Toronto ON, M5S 3H4, Canada

¹⁴Department of Physics and Astronomy, Tufts University, Medford, MA 02155, USA

Received 2025 March 18; revised 2025 June 1; accepted 2025 June 2; published 2025 July 14

## Abstract

We present photometric and spectroscopic observations of supernova SN 2023ixf covering from day 1 to 442 days after explosion. SN 2023ixf reached a peak V-band absolute magnitude of −18.2 ± 0.07, and light curves show that it has a relatively short “plateau” phase (~72 days), and can be classified either as an SN IIL or a transitional event between IIP and IIL. Early-time spectra of SN 2023ixf exhibit strong, very narrow emission lines from ionized circumstellar matter (CSM), possibly indicating a Type II classification. But these flash/shock-ionization emission features faded after the first week and the spectrum evolved in a manner similar to that of typical Type II SNe, unlike the case of most genuine SNe II in which the ejecta interact with CSM for an extended period of time and develop intermediate-width emission lines. We compare observed spectra of SN 2023ixf with various model spectra to understand the physics behind SN 2023ixf. Our nebular spectra (between 200 and 400 days) match best with the model spectra from a 15 $M_{\odot}$ progenitor that experienced enhanced mass loss a few years before explosion. A last-stage mass-loss rate of $\dot{M} = 0.01~M_{\odot}~\text{yr}^{-1}$ from the r1w6 model matches best with the early-time spectra, higher than $\dot{M} \approx 2.4 \times 10^{-3}~M_{\odot}~\text{yr}^{-1}$ derived from the ionized $H\alpha$ luminosity at 1.58 days. We also use SN 2023ixf as a distance indicator and fit the light curves to derive the Hubble constant by adding SN 2023ixf to the existing sample; we obtain $H_0 = 73.1_{-3.50}^{+3.68}~\text{km s}^{-1}~\text{Mpc}^{-1}$, consistent with the results from SNe Ia and many other independent methods.

Unified Astronomy Thesaurus concepts: Supernovae (1668); Type II supernovae (1731)

Materials only available in the online version of record: data behind figure

## 1. Introduction

Type II supernovae (hereafter SNe II) are characterized by prominent hydrogen Balmer lines in their optical spectra (e.g., A. V. Filippenko 1997). Thanks to progenitor detections (e.g., S. D. Van Dyk et al. 2003; S. J. Smartt 2009, 2015; S. D. Van Dyk 2017), it is commonly accepted that SNe II are explosions of massive stars ($\gtrsim 8~M_{\odot}$) at the end of their lives. While historically, SNe II were subgrouped based on their light-curve morphologies as SNe IIP with a long-duration

plateau and SNe IIL with faster linear (in magnitude) decline rates, recent studies suggest that the light curves actually form a more continuous distribution (e.g., J. P. Anderson et al. 2014; N. E. Sanders et al. 2015; L. Galbany et al. 2016; S. Valenti et al. 2016; T. de Jaeger et al. 2019). The duration of the plateau shows a continuous distribution ranging from 40 to 120 days (J. P. Anderson et al. 2014; D. Hiramatsu et al. 2021), with SNe IIP generally having a longer plateau; SNe IIL tend to have a shorter plateau because significant pre-explosion mass loss results in a partially stripped H-rich envelope. D. Hiramatsu et al. (2021) even argued that the continuity may be extended to SNe IIb, forming an SN IIP–IIL–IIb continuous distribution; SNe II with a very short plateau (40–70 days) and luminous optical peak ($\lesssim -18.4$ mag) are transitional between

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SNe IIL and IIb, with small H-rich envelope mass, high progenitor zero-age main sequence mass, and estimates of dense circumstellar matter (CSM). L. Dessart et al. (2024) also obtain a similar landscape of SNe IIP, IIL, IIb, and Ib by invoking binary progenitor massive stars of varying initial separation. Unlike mass loss in stellar winds, mass exchange in interacting binaries is a robust, powerful, mass-independent mass-loss mechanism that can produce a variety of partially or even fully stripped pre-SN stars.

However, even if SNe II constitute a uniform physical family, a great variety of objects is seen in terms of their photometric and spectroscopic properties. This diversity results from differences among progenitors and explosion mechanisms (e.g., mass of hydrogen envelope, radius, metallicity). For example, SNe II whose optical spectra exhibit long-lasting, intermediate-width emission lines that indicate expansion speeds of a few hundred to ~1000 km s⁻¹, but little or no PCygni absorption component, have been dubbed SNe IIn by E. M. Schlegel (1990); the “n” indicates the presence of the relatively “narrow” component. Examples were studied by A. V. Filippenko (1989, 1991a, 1991b), R. A. Stathakis & E. M. Sadler (1991), and others thereafter. The SN ejecta in these objects are likely interacting with CSM. However, C. L. Ransome et al. (2021) examine a wide variety of objects proposed to be SNe IIn, concluding that many of them do not even actually correspond to the terminal explosions of stars and are thus “SN impostors,” a term coined by S. D. Van Dyk et al. (2000) and subsequently used by many other authors.

In some cases, SNe IIn also show very narrow emission lines from gas having speeds <100 km s⁻¹; these are probably produced by flash or photoionization of CSM soon after shock breakout. Well-known examples include (among many others) SN 1998S (D. C. Leonard et al. 2000), SN 2013fs (O. Yaron et al. 2017), SN 2013cu (A. Gal-Yam et al. 2014), SN 2014G (G. Terreran et al. 2016), SN 2017ahn (L. Tartaglia et al. 2021), SN 2020pni (G. Terreran et al. 2022), and SN 2020tlf (W. V. Jacobson-Galán et al. 2022). These very narrow emission lines often disappear within a few days or a week, and the SN thereafter evolves in a manner more similar to that of typical Type II SNe (R. J. Bruch et al. 2023; W. V. Jacobson-Galán et al. 2024). Such objects, with only fleeting narrow emission lines and no long-lasting intermediate-width component, do not necessarily qualify as genuine SNe IIn, although opinions regarding this differ among authors.

Turning now to one aspect of their cosmological significance, SNe II have been established as useful independent extragalactic distance indicators (e.g., M. Hamuy & P. A. Pinto 2002; T. de Jaeger et al. 2020), especially for determining the Hubble–Lemaître constant (see T. de Jaeger & L. Galbany 2023 for a review) and addressing the current tension between early- and late-Universe measurements (A. G. Riess et al. 2022).

Owing to its exceptionally early detection in a nearby galaxy, SN 2023ixf provides a unique opportunity to explore and understand the zoo of SNe II; photometry and spectra were acquired starting <1 day after the explosion. Moreover, SN 2023ixf allows us to test the accuracy of different extragalactic distance methods found in the literature.

This paper is organized as follows. Section 2 contains a description of the observations and the data. In Section 3, we discuss the evolution of the light curves and color curves,

while in Section 4, we present the analysis of the spectroscopic data. Section 5 compares different methods to measure distances, and we conclude with a summary in Section 6.

## 2. Observations and Data Reduction

SN 2023ixf was discovered by Koichi Itagaki on 2023 May 19 at 17:27:15 (UTC dates are used throughout this paper) at a Clear-band magnitude of 14.9 (K. Itagaki 2023), in the nearby galaxy Messier 101 (M101, also known as NGC 5457 and informally as the “Pinwheel Galaxy”). Since M101 is a famous face-on spiral galaxy, it is regularly imaged by amateur astronomers throughout the world. After Itagaki reported the discovery of SN 2023ixf, many groups including amateurs and professionals reported prediscovery detections (V. Chufarin et al. 2023; A. V. Filippenko et al. 2023; M. D. Fulton et al. 2023; N. Hamann 2023; B. Koltenbah 2023; S. Limeburner 2023; Y. Mao et al. 2023; D. A. Perley & I. Irani 2023; O. Yaron et al. 2023; K. Zhang et al. 2023); see, in particular, the early-time unfiltered light curve of L. A. Sgro et al. (2023) obtained with Unistellar 11.4 cm telescopes and the multiband photometry of G. Li et al. (2024). These prediscovery observations range from a few minutes to a few days before Itagaki’s first report, and cover well before and after the explosion time of SN 2023ixf; thus, the explosion time can be precisely estimated (D. Hiramatsu et al. 2023; G. Hosseinzadeh et al. 2023). In this paper, we adopt the first-light time of SN 2023ixf to be MJD = 60082.788$_{-0.05}^{+0.02}$ estimated by G. Li et al. (2024).

Besides SN 2023ixf, M101 has also hosted a few other historical SNe, including SN 1909A (type unknown), SN 1951H (type unknown), SN 1970G (Type II), and SN 2011fe (Type Ia). Among these, the most important is SN Ia 2011fe since it is one of the best-observed SNe and can be used for precise distance (D) estimation. J. Vinkó et al. (2012) fit the multiband light curves of SN 2011fe using both the MLCS2k2 and SALT2 methods, deriving $D = 6.95 \pm 0.23$ Mpc (distance modulus $\mu = 29.21 \pm 0.07$ mag) from MLCS2k2 and $D = 6.46 \pm 0.21$ Mpc ($\mu = 29.05 \pm 0.07$ mag) from SALT2.

Given the proximity of M101, there are many other methods for measuring its distance. The most common and reliable one uses Cepheid variables. Over the past quarter century, several groups have independently measured the distance to M101 using Cepheids (W. L. Freedman et al. 2001; A. Saha et al. 2006; B. J. Shappee & K. Z. Stanek 2011). A. G. Riess et al. (2022) reported the latest Cepheid distance of $6.85 \pm 0.15$ Mpc ($\mu = 29.194 \pm 0.039$ mag). The optical tip of the red giant branch (TRGB) method is another option for distance measurement; R. L. Beaton et al. (2019) found $D = 6.52 \pm 0.12_{\text{stat}} \pm 0.15_{\text{sys}}$ ($\mu = 29.07 \pm 0.04_{\text{stat}} \pm 0.05_{\text{sys}}$ mag) using the latest TRGB method. Here for our absolute-magnitude analysis, we adopt $\mu = 29.194 \pm 0.039$ mag from A. G. Riess et al. (2022).

### 2.1. Photometry

We performed follow-up multiband observations of SN 2023ixf with both the 0.76 m Katzman Automatic Imaging Telescope (KAIT) and the 1 m Nickel telescope as part of the Lick Observatory Supernova Search (LOSS; A. V. Filippenko et al. 2001). *B*, *V*, *R*, and *I* images of SN 2023ixf were obtained with both telescopes; in addition, Clear-band (close to the *R* band; see W. Li et al. 2003) images were obtained with KAIT.
For the second and third nights, we also performed high-cadence observations with KAIT, rotating between the *B*, *V*, *R*, *I*, and Clear bands, with each exposure time being 60 s (plus an overhead of ~30 s), which gave a cadence of about 9 minutes for each band. (During the first night, unfortunately, winds at Lick Observatory were too strong for KAIT to be used; we had intended to conduct high-cadence photometry.)

All images were reduced using a custom pipeline<fnref n="15" /> detailed by B. E. Stahl et al. (2019). Point-spread-function photometry was obtained using DAOPHOT (P. B. Stetson 1987) from the IDL Astronomy Users Library.<fnref n="16" /> Owing to the small field of view of our images, only one reference star was available for calibration, namely star “m” from A. Henden et al. (2012, see their Figure 1). The A. U. Landolt (1992) magnitudes of star “m” were transformed to the KAIT/Nickel natural system before calibration. Apparent magnitudes were all measured in the KAIT4/Nickel2 natural system, and the final results were transformed to the standard system using the local calibrator and color terms for KAIT4 and Nickel2 (see B. E. Stahl et al. 2019).

## 2.2. Spectroscopy

We acquired 42 spectra (see Table 1 for a full log) of SN 2023ixf, generally with the Kast double spectrograph (J. Miller & R. Stone 1994) mounted on the Shane 3 m telescope at Lick Observatory. Most observations utilized the 2'' wide slit, 600/4310 grism, and 300/7500 grating (with a few exceptions). This instrument configuration has a combined wavelength range of ~3500–10500 Å and a spectral resolving power of $R \equiv \lambda/\Delta\lambda \approx 800$. To minimize slit losses caused by atmospheric dispersion (A. V. Filippenko 1982), the long slit was oriented at or near the parallactic angle. The data were reduced following standard techniques for CCD processing and spectrum extraction (J. M. Silverman et al. 2012) with IRAF (D. Tody 1986) routines and custom Python and IDL codes.<fnref n="17" /> Low-order polynomial fits to comparison-lamp spectra were used to calibrate the wavelength scale, and small adjustments derived from night-sky emission lines in the target frames were applied. The spectra were flux-calibrated using observations of appropriate spectrophotometric standard stars observed on the same night, at similar air masses, and with an identical instrument configuration. Among these spectra, some from the early phases have already been published by E. A. Zimmerman et al. (2024) and W. V. Jacobson-Galán et al. (2023; see Table 1 for details), but they are reused here for our more detailed analysis.

## 3. Light-curve Analysis

Figure 1 displays the full light curves of SN 2023ixf from our observations. The brightness of SN 2023ixf quickly rises to a peak in the first few days, especially in bluer bands, while the redder bands reach maximum brightness a bit later. There are noticeable fluctuations around the time of peak brightness in almost all bands, and the redder bands (*R* and *I*) actually peaked on the bumps. These early-time fluctuations are likely caused by CSM interaction, which was also detected in the early-time spectra (see Section 4). We use a low-order polynomial to fit the multiband light curves and determine

peak times of MJD = 60087.9, 60088.5, 60094.5, and 60097.3 (with $1\sigma$ uncertainties of 0.5 day) in the *B*, *V*, *R*, and *I* bands (respectively), as well as respective apparent peak magnitudes of 11.08, 11.11, 11.05, and 10.88 (with $1\sigma$ uncertainties of 0.05 mag). By adopting a first-light time of MJD = 60082.788$^{+0.02}_{-0.05}$ following G. Li et al. (2024), this gives rise times of $t_{r,B} = 5.1$ days, $t_{r,V} = 5.7$ days, $t_{r,R} = 11.7$ days, and $t_{r,I} = 14.5$ days (with $1\sigma$ uncertainties of ~0.5 day). Our rise-time estimates are consistent with the results given by W. V. Jacobson-Galán et al. (2023) in the *B* and *V* bands, but our *R* and *I* rise times are longer than their values, likely because our peak-time measurements in *R* and *I* landed on the bumps at later times. We also fit the $s1$ and $s2$ parameters to the *V* light curve (see Figure 1) following J. P. Anderson et al. (2014), finding $s1 = 3.1 \pm 0.1$ and $s2 = 1.7 \pm 0.1$ mag. This $s2$ value clearly puts SN 2023ixf in the fast-decliner (IIL) SN subclass instead of Type II plateau (IIP), at least according to their classification scheme.

To better compare with other SNe II, we plot the absolute V-band light curve in Figure 2 along with that of other SNe. To derive the absolute magnitude, we adopt a distance modulus $\mu = 29.194 \pm 0.039$ mag from A. G. Riess et al. (2022) and correct for both the Milky Way extinction of $E(B-V) = 0.008$ mag (E. F. Schlafly & D. P. Finkbeiner 2011) and the host-galaxy extinction of $E(B-V) = 0.032$ mag following S. D. Van Dyk et al. (2024). SN 2023ixf reached a peak V absolute magnitude of $-18.2 \pm 0.07$, at the bright end of SNe II (see also R. S. Teja et al. 2023) as shown in Figure 2, where a well-observed sample of SNe II from T. de Jaeger et al. (2019) is plotted for comparison (black lines). We also illustrate the typical Type IIP SN 2017eaw (S. D. Van Dyk et al. 2019) as well as the Type IIL SN 2006ai (D. Hiramatsu et al. 2021).

It is interesting to note that SN 2006ai not only has a similar peak absolute magnitude to SN 2023ixf, but also a relatively short plateau phase. Here we use the optically thick phase duration (OPTd; the time between the explosion and the plateau drop), following J. P. Anderson et al. (2014), to indicate the plateau phase, which is 72 days for SN 2023ixf and 63 days for SN 2006ai as given by D. Hiramatsu et al. (2021). SN 2006ai was argued to show its transitional nature between SN IIL and IIb by D. Hiramatsu et al. (2021). SN 2023ixf, with a duration ~10 days longer than SN 2006ai, could be more appropriately classified as a transitional event between SNe IIP and IIL, even though, according to the $s2$ value by J. P. Anderson et al. (2014), SN 2023ixf falls in the SN IIL class. In any case, both SN 2023ixf and SN 2006ai show a short plateau phase and thus clearly differ from the typical Type IIP SN 2017eaw, which lasts ~100 days before the light curve falls off the plateau.

A shorter plateau phase usually indicates an H-rich envelope of lower mass in the progenitor before explosion (D. J. Hillier & L. Dessart 2019). Model light curves for red supergiant (RSG) explosions with different H-rich envelope masses have been presented in a number of studies, including those by V. Morozova et al. (2015) and D. J. Hillier & L. Dessart (2019). Model “x5p0” from the latter matches closely the overall light curve of SN 2023ixf apart from the underestimate in early-time brightness (see Figure 2). This arises because model “x5p0” is a bare explosion (i.e., explosion within a vacuum), thus supporting the notion that the early-time luminosity boost stems from ejecta-CSM interaction.

<footnote n="15">¹⁵ https://github.com/benstahl92/LOSSPhotPypeline</footnote>

<footnote n="16">¹⁶ http://idlastro.gsfc.nasa.gov/</footnote>

<footnote n="17">¹⁷ https://github.com/ishivvers/TheKastShiv</footnote>
<table><caption>**Table 1**
Log of Lick/Kast Spectroscopic Observations</caption><thead><tr><th>Time (UTC)</th><th>MJD</th><th>Phase<sup>a</sup> (days)</th><th>Grism</th><th>Grating</th><th>Exp. Time (s)</th><th>Slit Width (arcsec)</th><th>Resolution (Å)</th></tr></thead><tbody><tr><td>20230520.183<sup>b</sup></td><td>60084.183</td><td>1.39</td><td>600/4310</td><td>300/7500</td><td>600/600</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20230520.196<sup>b</sup></td><td>60084.196</td><td>1.41</td><td>600/4310</td><td>600/5000</td><td>1200/1200</td><td>1.0</td><td>2.4/3.0</td></tr><tr><td>20230520.367<sup>b</sup></td><td>60084.367</td><td>1.58</td><td>600/4310</td><td>1200/5000</td><td>3660/3600</td><td>1.0</td><td>2.4/1.5</td></tr><tr><td>20230520.476<sup>b</sup></td><td>60084.476</td><td>1.69</td><td>600/4310</td><td>600/5000</td><td>1200/1200</td><td>1.0</td><td>2.4/3.0</td></tr><tr><td>20230520.488<sup>b</sup></td><td>60084.488</td><td>1.7</td><td>600/4310</td><td>300/7500</td><td>600/600</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20230521.19<sup>c</sup></td><td>60085.19</td><td>2.4</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>20230522.202<sup>d</sup></td><td>60086.202</td><td>3.4</td><td>600/4310</td><td>600/7500</td><td>600/600</td><td>2.0</td><td>4.7/6.0</td></tr><tr><td>20230523.231<sup>d</sup></td><td>60087.231</td><td>4.4</td><td>600/4310</td><td>600/7500</td><td>600/600</td><td>2.0</td><td>4.7/6.0</td></tr><tr><td>20230524.313<sup>d</sup></td><td>60088.313</td><td>5.5</td><td>600/4310</td><td>600/7500</td><td>360/360</td><td>1.0</td><td>2.4/3.0</td></tr><tr><td>20230525.197<sup>d</sup></td><td>60089.197</td><td>6.4</td><td>600/4310</td><td>600/7500</td><td>400/400</td><td>1.0</td><td>2.4/3.0</td></tr><tr><td>20230527.210<sup>d</sup></td><td>60091.210</td><td>8.4</td><td>600/4310</td><td>600/7500</td><td>180/180</td><td>1.0</td><td>2.4/3.0</td></tr><tr><td>20230528.189</td><td>60092.189</td><td>9.4</td><td>600/4310</td><td>300/7500</td><td>200/200</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20230528.196</td><td>60092.196</td><td>9.4</td><td>600/4310</td><td>600/5000</td><td>500/500</td><td>1.0</td><td>2.4/3.0</td></tr><tr><td>20230528.435</td><td>60092.435</td><td>9.6</td><td>600/4310</td><td>300/7500</td><td>300/300</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20230528.445</td><td>60092.445</td><td>9.7</td><td>600/4310</td><td>600/5000</td><td>1200/1200</td><td>1.0</td><td>2.4/3.0</td></tr><tr><td>20230530.247<sup>d</sup></td><td>60094.247</td><td>11.5</td><td>600/4310</td><td>600/7500</td><td>360/360</td><td>1.0</td><td>2.4/3.0</td></tr><tr><td>20230531.193<sup>d</sup></td><td>60095.193</td><td>12.4</td><td>600/4310</td><td>600/7500</td><td>210/210</td><td>1.0</td><td>2.4/3.0</td></tr><tr><td>20230601.194</td><td>60096.194</td><td>13.4</td><td>600/4310</td><td>600/7500</td><td>300/300</td><td>1.0</td><td>2.4/3.0</td></tr><tr><td>20230602.258<sup>d</sup></td><td>60097.258</td><td>14.5</td><td>600/4310</td><td>300/7500</td><td>200/200</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20230610.262</td><td>60105.262</td><td>22.5</td><td>600/4310</td><td>300/7500</td><td>300/300</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20230619.246</td><td>60114.246</td><td>31.5</td><td>600/4310</td><td>300/7500</td><td>600/600</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20230622.266</td><td>60117.266</td><td>34.5</td><td>600/4310</td><td>300/7500</td><td>200/200</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20230626.242</td><td>60121.242</td><td>38.5</td><td>600/4310</td><td>300/7500</td><td>200/200</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20230710.233</td><td>60135.233</td><td>52.4</td><td>600/4310</td><td>300/7500</td><td>200/200</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20230719.319</td><td>60144.319</td><td>61.5</td><td>600/4310</td><td>300/7500</td><td>200/200</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20230811.246</td><td>60167.246</td><td>84.5</td><td>600/4310</td><td>300/7500</td><td>1200/1200</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20230824.232</td><td>60180.232</td><td>97.4</td><td>600/4310</td><td>300/7500</td><td>600/600</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20230908.158</td><td>60195.158</td><td>112.4</td><td>600/4310</td><td>300/7500</td><td>461/600</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20230916.144</td><td>60203.144</td><td>120.4</td><td>...</td><td>300/7500</td><td>...</td><td>2.0</td><td>...</td></tr><tr><td>20231013.099</td><td>60230.099</td><td>147.3</td><td>600/4310</td><td>300/7500</td><td>300/300</td><td>0.5</td><td>1.2/3.0</td></tr><tr><td>20231109.565</td><td>60257.565</td><td>174.8</td><td>600/4310</td><td>300/7500</td><td>1200/1200</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20231110.561</td><td>60258.561</td><td>175.8</td><td>600/4310</td><td>300/7500</td><td>1200/1200</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20231212.534</td><td>60290.534</td><td>207.7</td><td>600/4310</td><td>600/5000</td><td>2460/2400</td><td>1.0</td><td>2.4/3.0</td></tr><tr><td>20240112.504</td><td>60321.504</td><td>238.7</td><td>600/4310</td><td>600/5000</td><td>3060/3000</td><td>1.0</td><td>2.4/3.0</td></tr><tr><td>20240118.561</td><td>60327.561</td><td>244.8</td><td>600/4310</td><td>300/7500</td><td>2160/1400</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20240207.668<sup>e</sup></td><td>60347.668</td><td>264.9</td><td>600/4000</td><td>400/8500</td><td>250/250</td><td>1.0</td><td>4.7/8.6</td></tr><tr><td>20240313.658<sup>e</sup></td><td>60382.658</td><td>299.9</td><td>600/4000</td><td>400/8500</td><td>180/180</td><td>1.0</td><td>4.7/8.6</td></tr><tr><td>20240411.314</td><td>60411.314</td><td>328.5</td><td>600/4310</td><td>300/7500</td><td>2760/2700</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20240515.406</td><td>60445.406</td><td>362.6</td><td>600/4310</td><td>300/7500</td><td>2760/2700</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20240531.449</td><td>60461.449</td><td>378.7</td><td>600/4310</td><td>300/7500</td><td>3060/3000</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20240629.352</td><td>60490.352</td><td>407.6</td><td>600/4310</td><td>300/7500</td><td>3060/3000</td><td>2.0</td><td>4.7/11.9</td></tr><tr><td>20240802.276</td><td>60524.276</td><td>441.5</td><td>600/4310</td><td>300/7500</td><td>3060/3000</td><td>2.0</td><td>4.7/11.9</td></tr></tbody></table>

**Notes.**

<sup>a</sup> Relative to first-light time of MJD = 60082.788.

<sup>b</sup> Already published by E. A. Zimmerman et al. (2024).

<sup>c</sup> Adopted from W. V. Jacobson-Galán et al. (2023) for analysis.

<sup>d</sup> Already published by W. V. Jacobson-Galán et al. (2023), but with our independent reduction.

<sup>e</sup> Taken with 10 m Keck telescope + LRIS.

The B-V color evolution of SN 2023ixf and the comparison sample is also shown in Figure 2 in the middle panel. SN 2023ixf, SN 2006ai, SN 2017eaw, and the rest of the sample all have very similar B-V evolution, with a blue color at the beginning evolving to a redder color. We plot the uvw2-R evolution in the bottom panel, where the uvw2-band data were adopted from E. A. Zimmerman et al. (2024); this shows that SN 2023ixf reached its bluest UV/optical color ~5 days after explosion.

Figure 3 shows our high-cadence light curves from Night 2 (top panel) and Night 3 (bottom panel). Such data can be used to search for rapid variability in SN light curves, as done for SN 2014J by A. Z. Bonanos & P. Boumis (2016), who performed observations with a cadence of 2 minutes in B and V, finding evidence for rapid variability at the 0.02–0.05 mag level on timescales of 15–60 minutes. For SN 2023ixf, we found no evidence of similar rapid variability in the second and third nights, since our high-cadence light curves (as shown
<figure><img src="image_3.png" /><figcaption>**Figure 1.** Left: LOSS multiband light curves of SN 2023ixf, color-coded for different bands. Right: zoomed-in light curves at early phases. The s1 and s2 parameters are shown for the V-band fitting following J. P. Anderson et al. (2014). With an s2 value of 1.7 mag, SN 2023ixf is clearly in the fast-decliner (IIL) subclass instead of a typical SN IIP. The photometry is available as the data behind this figure in machine-readable format.</figcaption></figure>

(The data used to create this figure are available.)

in Figure 3) are quite smooth, and also because our cadence is much lower. However, our light curves clearly show that the SN brightened steadily during both nights. Here we define $s'$ as the daily brightening rate; we find $s' = 1.0 \pm 0.1 \text{ mag day}^{-1}$ for all bands in Night 2 and $s' = 0.4 \pm 0.1 \text{ mag day}^{-1}$ for all bands in Night 3, so the SN brightened more quickly during the second night than the third. A higher cadence with fewer filters may help reveal rapid variability (if it exists) in future observations of bright SNe.

## 4. Spectral Analysis

Figure 4 shows our full set of Kast spectra (a total of 42; see Table 1 for a full log) of SN 2023ixf until 442 days after explosion. Our first spectrum was taken at MJD = 60084.183, merely 1.39 days after explosion. We obtained nearly daily Kast spectra of SN 2023ixf for the first two weeks after explosion. In addition, for the first-night observations, we obtained five spectra ranging from 1.4 days to 1.7 days with different grating resolution with the same telescope (note that K. A. Bostroem et al. 2023 also obtained multiple spectra in the same night, but with different telescopes). These five spectra from the first-night observations were already shown by E. A. Zimmerman et al. (2024); also, the 3.4 days, 4.4 days, 5.5 days, 6.4 days, and 8.4 days spectra were presented by W. V. Jacobson-Galán et al. (2023), but we include them here for a more complete and detailed analysis. Augmenting the spectral set, we plot the classification spectrum of SN 2023ixf at 1.14 days from the Liverpool Telescope; this is the earliest reported spectrum of SN 2023ixf (D. A. Perley et al. 2023). These high-quality spectra allow us to perform thorough analysis of the spectral evolution of SN 2023ixf.

Distinct spectral evolution was seen from day 1.4 to day 442 as shown in Figure 4. We can classify the evolution into four

main phases, following the suggestion given by L. Dessart et al. (2016) for SN 1998S. During the first few days, the most distinct features are the strong but narrow emission lines from the ionized CSM; we call this the CSM-dominated phase. These emission features weakened after the first week, so the spectra appear to be blue and quasi-featureless in the second week; this is the cold-dense-shell (CDS)-dominated phase, during which the only distinct features are from hydrogen, though weak and broad P Cygni features started to emerge at the end of the second week. After ~3 weeks, during the ejecta-dominated phase, the spectra are similar to those of other SNe II with strong P Cygni lines. After ~100 days, the spectra enter the nebular phase and forbidden lines start to form at late times. Detailed descriptions of each phase are given in the following sections.

### 4.1. CSM-dominated Phase

Figure 5 presents spectra of SN 2023ixf in the first five days during the CSM-dominated phase. During this phase, there are many both strong and weak narrow emission lines apparently caused by the ionized CSM, also reported by other groups (K. A. Bostroem et al. 2023; D. Hiramatsu et al. 2023; G. Hosseinzadeh et al. 2023; W. V. Jacobson-Galán et al. 2023; R. S. Teja et al. 2023; M. Yamanaka et al. 2023). These lines are also accompanied by broad components that are caused by electron scattering. We identified many species and found that most are from H and He.

The strongest emission in the CSM-dominated phase is the H I Balmer series including H$\alpha$, H$\beta$, H$\gamma$, H$\delta$, H$\epsilon$, H$\zeta$, and H$\eta$, owing to the large abundance of hydrogen. As expected, H$\alpha$ is the strongest line, with other Balmer lines gradually weaker at higher energy states (bluer wavelengths). In addition, the H I Paschen series also appears in our spectra at 1.41 days,
<figure><img src="image_4.png" /></figure>

**Figure 2.** Top panel: absolute V-band light curve of SN 2023ixf (red) compared with the typical Type IIP SN 2017eaw (green), the Type IIL SN 2006ai, the model "x5p0" (purple) from D. J. Hillier & L. Dessart (2019) and a well-observed sample of SNe II (black) from T. de Jaeger et al. (2019). SN 2023ixf lies at the bright end of the Type II sample, similar to SN 2006ai in absolute peak V-band magnitude, and exhibits a relatively short "plateau" phase (less than ~70 days). Middle panel: the B–V color evolution of SN 2023ixf and the comparison sample; all show very similar B–V curves, with blue colors at the beginning evolving to redder colors. Bottom panel: the uvw2–R color evolution (with a zoomed-in panel) of SN 2023ixf.

1.70 days, and 2.4 days. Although Pa$\alpha$ λ12157 is outside our wavelength coverage, Pa$\beta$, Pa$\gamma$, Pa$\delta$, Pa$\epsilon$, Pa$\zeta$, and Pa$\eta$ are clearly detected. However, the H I Paschen series is much weaker than the H I Balmer series. The H I Balmer and Paschen series are marked with red vertical lines in Figure 5.

Helium also exhibits some strong emission lines. In fact, He II λ4686 (at an energy state of *n* = 3) is the second-strongest feature during the CSM-dominated phase. Moreover, a few other He II lines are detected from energy states *n* = 4, 5, and 6, marked as green vertical lines in Figure 5. A list of He II wavelengths from different energy states is given in Table 2, taken from the Atomic Data as part of CMFGEN<fnref n="18" /> (D. J. Hillier & L. Dessart 2012). Such complete He II series in SN spectra are presented here for the first time (to our knowledge), although some lines at higher energy states are too weak to be clearly detected. Spectra with higher signal-to-noise ratio (S/N) could potentially reveal these weak lines in future SNe showing ionization features. In addition to He II, there are also a few He I lines detected; though much weaker, they are still quite obvious in our first-night spectra, including He I λ5876, λ6678, and λ7065.

Besides hydrogen and helium, a few C and N features (marked as cyan and blue vertical lines) were also identified in the early-time spectra. C IV λλ5801, 5812 is not only detected, but even distinguished separately in our higher-resolution

<figure><img src="image_5.png" /></figure>

**Figure 3.** High-cadence KAIT light curves during Night 2 (top panel) and Night 3 (bottom panel). The smooth light curves show no evidence of rapid variability in either night. The SN brightened steadily, with rates of ~1.0 and ~0.4 mag per day during Nights 2 and 3, respectively. The photometry is available as the data behind this figure in machine-readable format.
(The data used to create this figure are available.)

spectra on days 1.41, 1.58, and 1.69 (see also N. Smith et al. 2023 and D. Dickinson et al. 2025 with high-resolution spectroscopy). The detection of C IV λ4658 is uncertain as it is blended with other lines. No lower ionization state C III lines are obviously detected; C III λλ4647, 4650 might be present, but it is blended with C IV λ4658. N IV lines are also clearly detected with N IV λ4058 and λλ7109, 7123. The lower ionization state N III λλ4634, 4641 is visible, but N III λ4687 is blended with He II λ4686 and thus uncertain. Both N IV λλ7109, 7123 and N III λλ4634, 4641 are also identified in our higher-resolution spectra on days 1.41, 1.58, and 1.69, similar to C IV λλ5801, 5812 (N. Smith et al. 2023 also show high-resolution spectra).

Interestingly, thanks to the total of five spectra taken during the first night of observations, we find that a few of the above-mentioned lines evolve quickly, showing intraday changes (see also K. A. Bostroem et al. 2023). For example, all three detected He I lines (λ5876, λ6678, and λ7065) weakened and almost disappeared from days 1.41 to 1.70, within 8 hr. To better illustrate this, in Figure 6 we plot the three He I regions centered on C IV λλ5801, 5812, H$\alpha$, and N IV λ7115. By day 2.4 after explosion, all three He I lines disappeared completely. In addition, the N III λλ4634, 4641 and C III λλ4647, 4650 lines (blended) weakened quickly in the first two days, as

<footnote n="18">¹⁸ http://kookaburra.phyast.pitt.edu/hillier/web/CMFGEN.htm</footnote>
<figure><img src="image_6.png" /><figcaption>**Figure 4.** Lick/Kast observations of SN 2023ixf, with a total of 42 spectra covering from day 1.4 to day 442 after explosion. The 1.14 days spectrum (shown in red) was taken from D. A. Perley et al. (2023) for completeness. Distinct spectral evolution is seen. The spectra in the first week present strong but narrow emission lines from the ionized CSM (CSM-dominated phase), which subsequently weakened; the spectra appear to be featureless in the second week (CDS-dominated phase), though weak broad P Cygni features start to emerge. After ~3 weeks, the spectra are similar to those of other SNe II with strong P Cygni lines (ejecta-dominated phase). The spectra enter the nebular phase after ~100 days and forbidden lines start to form at late times. These spectra are available as the data behind the figure in individual machine-readable tables in a .tar.gz package.
(The data used to create this figure are available.)</figcaption></figure>

shown in Figure 7 centered on He II λ4686. The profile is very complicated, with many lines blended in the blue wing of He II λ4686, including clear detections of C III λλ4647, 4650, N III λλ4634, 4641, and He II λ4542, and possibly C IV λ4658 and N V λλ4604, 4620. These lines boost the flux of the blue wing (while the red wing has a normal shape), making the He II profile very asymmetric. The emission blueward of He II λ4686 even strengthens over the first few days, as shown in the

right panel of Figure 7. Our highest-resolution spectrum at 1.58 days is able to distinguish both components of the N III λλ4634, 4641 doublet, which was even stronger in the 1.14 days spectrum from the Liverpool Telescope, comparable to He II λ4686, but weakened significantly to nearly undetected by day 1.70. Interestingly, there appears to be N V λλ4604, 4620 in the 2.4 days spectrum, but not earlier nor later. If real, this is additional evidence for intraday changes of the quickly evolving lines.

## 4.2. CDS-dominated Phase

Overall, almost all the emission features faded away after about a week (except for the H I Balmer series, mainly Hα), and the spectra became quasi-featureless with a blue continuum during the second week, indicating that the ejecta are CDS-dominated at this time, as shown in Figure 8. Because of this, one can try to fit the quasi-featureless spectra with a blackbody function to estimate the temperature of the ejecta, given the reliable relative flux calibration. To do so, we first correct the spectral extinction by adopting a total E (B − V) = 0.040 mag as mentioned above. We then exclude regions with obvious emission lines (we applied the same blackbody fitting to the spectra from the CSM-dominated phase, too). Finally, in order to keep consistent for spectra with different wavelength coverage at the red end, we restrict the fitting region to be below 8600 Å (the blue-end coverage is almost the same in all spectra). A blackbody function was applied to fit the spectra and the results are shown in Figure 9; the top panel indicates how well the blackbody function (blue) matches the selected data (red) by excluding the emission lines and the region beyond 8600 Å (black), and the bottom panel shows the evolution of the fitted temperature (blue circles) with time.

Our results demonstrate that the blackbody temperature rises quickly during the first few days, with ~16,000 K at 1.4 days and quickly reaching a peak around 24,000 K about 2.5 days after explosion. The temperature then gradually drops in the next two weeks, and by day 14 it is ~13,000 K. For comparison, we also plot the temperature fitting results from E. A. Zimmerman et al. (2024, green triangles), where they use the photometric fitting method (including UV photometry). Our results are generally consistent with theirs, though it appears that our fitted temperatures are slightly lower; however, they are higher than the fitted temperatures given by J. Zhang et al. (2023, see their Figure S3, middle panel). Considering that E. A. Zimmerman et al. (2024) have better wavelength coverage (including UV photometry) than our optical spectra, and the peak blackbody emission lies in the far-UV region, it is reasonable to consider that their fitting results better constrain the temperatures. Our spectral fitting may count as an independent check on the overall temperature evolution.

In order to better reveal the morphology of the weak lines in the quasi-featureless spectra, we adopt a normalization procedure following that of D. C. Leonard et al. (2000, see their Figure 5): the spectra are divided by the continuum, which was fitted with several-order splines by excluding obvious emission and absorption regions. The normalized spectra are shown in Figure 10.

During this phase, one interesting feature is that the Hα line starts to show a low-velocity P Cygni profile on day 8.4, as shown in Figure 10 (also see Figures 11 and 12). The velocity
<figure><img src="image_7.png" /><figcaption>**Figure 5.** Lick/Kast spectra of SN 2023ixf from the first five days after explosion during the CSM-dominated phase. Many species are detected and identified, including the H I Balmer and Paschen series marked with red vertical lines; the He II series from energy states of *n* = 3, 4, 5, and 6, as well as a few He I transitions marked with green vertical lines; and some C and N transitions marked with cyan and blue vertical lines. The red spectra were from model r1w6 at a nearby phase for comparison; see text for details.</figcaption></figure>

from the blue absorption minimum is <1000 km s⁻¹ and lasts for about a week, until day 14.5 (before a gap in our spectral coverage in the third week). The velocity is likely caused by radiative acceleration (L. Dessart 2024), and it is still unclear whether this mechanism could last for such a long time (see Section 4.5). However, there are no such low-velocity P Cygni features in the other H I Balmer lines beyond Hα (perhaps a weak signature in Hβ on day 9; see the middle panel in Figure 12). Instead, a broad P Cygni profile (ν > 5000 km s⁻¹) starts to emerge in Hβ, Hγ, and Hδ on day 8.4 (see next section for velocity measurements).

The Hα line at this phase does not exhibit a broad P Cygni profile, probably because the CDS at this phase is still optically thick. The H I Paschen series, however, remains as narrow emission until day 14.5. Unlike the H I Balmer lines, no low-velocity P Cygni nor broad P Cygni profiles are clearly seen in the H I Paschen series.

Interestingly, the He I lines that earlier disappeared (narrow emission lines) also start to be visible again, with broad P Cygni profiles around day 8, as shown in Figure 13. The He I λ5876 profile in the left panel shows obvious absorption at a velocity of ~7000 km s⁻¹, though the absorption in the He I

λ7065 profile (right panel) is not significant. Also, there appears to be another bluer absorption in He I λ5876 at 14.5 days with a velocity around 12,000 km s⁻¹, which seems a bit high but not impossible.

### 4.3. Ejecta-dominated Phase

Unfortunately no spectrum was obtained during the third week after explosion. After three weeks, the spectra taken at day 22.5 and thereafter show that the SN is ejecta-dominated; there are many broad P Cygni features that are seen in normal SNe II. Figure 14 shows the spectra from days 22.5 to 120.4. Broad P Cygni profiles of H I Balmer lines are very obvious, as well as those of a few other lines such as Fe II λλλλ4924, 5018, 5169.

We measure the velocity from the blue absorption minimum of the broad P Cygni profile, and the results are shown in Figure 15. The velocity evolution after three weeks is typical of all lines compared with other normal SNe II; they all decline gradually with a power-law shape, with the Hα velocity higher than that of other lines. There is a dip blueward of Hα absorption from days 22.5 to 38.5, likely due to Si II λ6355 at ~5500 km s⁻¹ (unlikely to be high-velocity Hα, which would
<figure><img src="image_8.png" /><figcaption>**Figure 6.** Zoomed-in regions of three detected He I lines, λ5876 (left), λ6678 (middle), and λ7065 (right), respectively centered on C IV λλ5801, 5812, Hα, and N IV λ7115. All three He I lines weakened and almost disappeared from day 1.41 to day 1.70, and by day 2.4 they vanished completely.</figcaption></figure>

<table><caption>**Table 2**

He II Air Wavelengths for Energy States *n* = 3, 4, 5, and 6</caption><thead><tr><th></th><th colspan="2">n = 3</th><th colspan="2">n = 4</th><th colspan="2">n = 5</th><th colspan="2">n = 6</th></tr><tr><th></th><th>n'</th><th>λ (Å)</th><th>n'</th><th>λ (Å)</th><th>n'</th><th>λ (Å)</th><th>n'</th><th>λ (Å)</th></tr></thead><tbody><tr><td>4</td><td>4685.70</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>5<sup>a</sup></td><td>3203.09</td><td>5</td><td>10123.59</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>...</td><td>...</td><td>6</td><td>6560.09</td><td>6<sup>a</sup></td><td>18636.75</td><td>...</td><td>...</td><td>...</td></tr><tr><td>...</td><td>...</td><td>7</td><td>5411.52</td><td>7<sup>a</sup></td><td>11626.41</td><td>7<sup>a</sup></td><td>30908.46</td><td>...</td></tr><tr><td>...</td><td>...</td><td>8</td><td>4859.32</td><td>8</td><td>9344.93</td><td>8<sup>a</sup></td><td>18743.33</td><td>...</td></tr><tr><td>...</td><td>...</td><td>9</td><td>4541.59</td><td>9</td><td>8236.79</td><td>9<sup>a</sup></td><td>14760.38</td><td>...</td></tr><tr><td>...</td><td>...</td><td>10</td><td>4338.68</td><td>10<sup>a</sup></td><td>7592.76</td><td>10<sup>a</sup></td><td>12812.83</td><td>...</td></tr><tr><td>...</td><td>...</td><td>11</td><td>4199.84</td><td>11</td><td>7177.53</td><td>11<sup>a</sup></td><td>11673.25</td><td>...</td></tr><tr><td>...</td><td>...</td><td>12</td><td>4100.05</td><td>12</td><td>6890.91</td><td>12<sup>a</sup></td><td>10933.62</td><td>...</td></tr><tr><td>...</td><td>...</td><td>13</td><td>4025.61</td><td>13</td><td>6683.21</td><td>13</td><td>10419.83</td><td>...</td></tr><tr><td>...</td><td>...</td><td>14<sup>a</sup></td><td>3968.44</td><td>14</td><td>6527.11</td><td>14</td><td>10045.27</td><td>...</td></tr><tr><td>...</td><td>...</td><td>15<sup>a</sup></td><td>3923.49</td><td>15</td><td>6406.39</td><td>15</td><td>9762.17</td><td>...</td></tr><tr><td>...</td><td>...</td><td>...</td><td>...</td><td>16</td><td>6310.86</td><td>16</td><td>9542.07</td><td>...</td></tr><tr><td>...</td><td>...</td><td>...</td><td>...</td><td>17<sup>a</sup></td><td>6233.83</td><td>17</td><td>9367.05</td><td>...</td></tr><tr><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>18</td><td>9225.25</td><td>...</td></tr><tr><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>19</td><td>9108.55</td><td>...</td></tr><tr><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>20<sup>a</sup></td><td>9011.23</td><td>...</td></tr><tr><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>21<sup>a</sup></td><td>8929.13</td><td>...</td></tr><tr><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>22<sup>a</sup></td><td>8859.17</td><td>...</td></tr><tr><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>23<sup>a</sup></td><td>8799.02</td><td>...</td></tr></tbody></table>

**Note.**

<sup>a</sup> Either not in our wavelength range or not obviously detected owing to low S/N.

be at ~15,000 km s⁻¹). Si II λ6355 is also predicted by the model (see model spectrum at 23.6 days in Figure 14). However, the velocity of Si II λ6355 (~5500 km s⁻¹) appears to be a bit low compared with Hα, so it is possible this is caused by some other line. The Fe II λλλ4924, 5018, 5169 lines also evolve in a manner similar to that of T. Faran et al. (2014) as shown in Figure 15, though they are slightly higher than the results of C. P. Gutiérrez et al. (2017).

However, it is interesting to see that the velocities of Hβ, Hγ, and Hδ (emerging earlier than Hα) at early times actually rise in the second week and reach their peak at ~20 days, ~8000 km s⁻¹. Such a rise in velocity at early times is very unusual in SNe II, though it is predicted by some models (see,

e.g., Figure 9 of L. Dessart & D. J. Hillier 2010). This is likely not related to the unshocked CSM because the velocities are too high. Instead, it is a projection effect caused by sphericity (L. Dessart & D. J. Hillier 2005, 2010).

### 4.4. Nebular Phase

After ~100 days, the spectra enter the nebular phase and forbidden lines start to form at late times. Although H I Balmer lines remain the strongest features (mostly still absorption, as well as absorption from Fe II λλλ4924, 5018, 5169), strong emission from forbidden lines also starts to emerge, such as [O I] λλ6300, 6364 and λ7774, Mg I] λ4571, [Ca II] λλ7291, 7323, and Ca II λλλ8498, 8542, 8662, as shown in Figure 16.

### 4.5. Models and Physics

In order to produce the strongly ionized features (A. Gal-Yam et al. 2014) seen in the early-time spectra of SN 2023ixf, a certain amount of dense CSM is required surrounding the progenitor star, which could extend the shock-breakout time to days (compared to a timescale of hours if without dense CSM, thus difficult to detect). As the radiation leaks from the shock, UV photons are produced from the cooling of the shock-heated CSM, which then ionize the CSM, forming these features. Such models were studied/simulated by L. Dessart et al. (2017, hereafter D17), where a set with different parameter spaces (radius, mass-loss rate, ejecta mass, etc.) were explored after explosion from an RSG progenitor. In fact, the progenitor of SN 2023ixf has been directly identified and confirmed to be an RSG (e.g., J. E. Jencson et al. 2023; C. D. Kilpatrick et al. 2023; Z. Niu et al. 2023; J. L. Pledger & M. M. Shara 2023; M. D. Soraisam et al. 2023; Y.-J. Qin et al. 2024; S. D. Van Dyk et al. 2024; D. Xiang et al. 2024), with an effective temperature, luminosity, and initial mass of ~3450 K, ~9.3 × 10⁴ L⊙, and 12–15 M⊙, respectively (S. D. Van Dyk et al. 2024). In the D17 models, r1w4 and r1w6 are the two showing strongly ionized features at early times, but r1w6 matches the evolution of SN 2023ixf better since the r1w4 model evolves too fast, so here we focus more on r1w6. For the r1w6 model, the following parameters were set: progenitor star radius R<sub>⋆</sub> = 501 R⊙ (i.e., ~3.5 × 10¹³ cm), mass-loss rate M̄ = 10⁻² M⊙ yr⁻¹, extended
<figure><img src="image_9.png" /><figcaption>**Figure 7.** Left: zoomed-in regions of He II λ4686 along with Hβ and Hγ centered on He II λ4686. The blue-wing profile of the He II λ4686 region is very complicated, with many lines blended together, including clear detections of C III λλ4647, 4650, N III λλ4634, 4641, and He II λ4542, as well as possible C IV λ4658 and N V λλ4604, 4620. These lines boost the flux of the blue wing, making the He II profile very asymmetric. Right: the same region but with spectra arbitrarily scaled to show that the emission blueward of He II λ4686 strengthens over the first few days.</figcaption></figure>

<figure><img src="image_10.png" /><figcaption>**Figure 8.** Lick/Kast spectra of SN 2023ixf from the second week after explosion in the CDS-dominated phase, during which they become featureless; almost all emission lines faded away after about a week, except Hα (other H I Balmer and Paschen series lines are quite weak and are marked as red vertical lines). The Hα line starts showing a low-velocity (less than 1000 km s⁻¹) P Cygni profile on day 8, and a broad P Cygni profile starts to emerge in the bluer region (Hβ, Hγ, Hδ, etc.) on day 9.4. The spectra shown in red are from the r1w6 model at a nearby phase for comparison; see text for details. Note that the drop in the model at the blue edge after 7 days corresponds to the Balmer edge, at which the model predicts erroneously a sharp jump.</figcaption></figure>

CSM mass $M_{\text{ext}} = 3.04 \times 10^{-2} M_{\odot}$ up to a distance of $R_{\text{CSM}} \approx 5 \times 10^{14}$ cm, and ejecta mass $M_{\text{ejecta}} = 12.52 M_{\odot}$ (see Table 1 and Section 3 of L. Dessart et al. 2017 for more details). A montage showing the evolution of the ejecta/CSM properties from the radiation hydrodynamical simulation of model r1w6 (L. Dessart et al. 2017) is plotted in Figure 17, where black, blue, and red lines respectively show the variation of velocity, temperature, and luminosity with radial distance at various epochs. The shaded area represents regions with $\tau_{\text{es}} > 2/3$.

The r1w6 model spectra<fnref n="19" /> at 1.5 days, 2.0 days, and 3.0 days are plotted in Figure 5 in the CSM-dominated phase and compared with the observed spectra. Overall, the model spectra match the observed spectra quite satisfactorily, including most of the major features such as the H I Balmer and Paschen series, all the He II series, C IV, N III, and N IV. However, there do exist some mismatches. (1) As mentioned in Section 4.1, there are three He I lines (λ5876, λ6678, and λ7065) detected in the first-night spectra (days 1.41–1.70), yet they are not predicted (at least not obvious) in the model spectra,<fnref n="20" /> and they evolve quickly, showing intraday changes and almost disappearing within 8 hr. The disappearance of the He I lines is likely because of the quick increase in the temperature during the second day after explosion. As shown in Figure 9, the temperature rises quickly in the first three days,

<footnote n="19">¹⁹ Those used in this paper are slightly updated compared with the original r1w6 model given by D17 in which only H, He, CNO, and Fe were included. The updated models presented here include all important metals up to nickel together with a more extended model atom for up to five ions per species (Fe I to Fe V) as well as number of levels etc.</footnote>

<footnote n="20">²⁰ The He I lines are predicted in a slight variant of the r1w6 model presented by W. V. Jacobson-Galán et al. (2023); see their Figure 2.</footnote>
<figure><img src="image_11.png" /><figcaption>**Figure 9.** Top: blackbody fitting to the observed spectra, where blue lines show the results of the blackbody function and red indicates the selected fitting region after excluding emission lines and beyond 8600 Å (black). Bottom: fitted evolution of temperature (blue circles) with time. The ejecta temperature rises quickly in the first few days and reaches a peak around 24,000 K at ~2.5 days after explosion. It then gradually drops in the next two weeks. Similar fitting results from E. A. Zimmerman et al. (2024) are also plotted for comparison (green triangles), showing consistent results.</figcaption></figure>

being ~16,000 K at 1.4 days and quickly reaching a peak around 24,000 K at ~2.5 days after explosion. The He I lines require lower ionization energy (temperature); thus, at $t < 1.7$ days, with lower temperatures, some He I lines formed and are detected in our spectra. As the temperature increases quickly and reaches its peak at ~2.5 days, most of the helium was ionized to He II, which requires a higher ionization energy (temperature) than He I, so He I almost disappears after two days, as predicted by the models too. (2) There appears to be $N\,V\,\lambda\lambda4604$, 4620 in the spectrum at 2.4 days, but not before nor after. If real, this is likely also because of the changes in temperature. Typically, owing to stratification of temperature/ionization in the CSM, the $N\,V$ forms in the deepest layers of the CSM (i.e., closer to the shock where temperature is higher) as $N\,V$ requires a very high ionization energy/temperature (there could also be a greater diversity in ionization if the CSM is asymmetric). At day 2.4, the temperature reached its peak, meeting the requirements to form $N\,V$ lines (but not before the peak), and the temperature also decreases after the peak, so the $N\,V$ disappeared thereafter. (3) The blue wing of He II $\lambda4686$ is stronger and more asymmetric than the red wing, differing from the model spectra, and were seen in some SNe (W. V. Jacobson-Galán et al. 2024). Although the blue wing

<figure><img src="image_12.png" /><figcaption>**Figure 10.** Spectra at the same epochs as in Figure 9, but normalized by dividing by the continuum. Red vertical lines mark the H I Balmer and Paschen series, and green vertical lines mark He I $\lambda5876$, $\lambda6678$, and $\lambda7065$. Red spectra are the same as given in Figure 8.</figcaption></figure>

<figure><img src="image_13.png" /><figcaption>**Figure 11.** Zoomed-in regions of $H\alpha$ profile from days 6.4 to 11.5. The $H\alpha$ line starts to show a low-velocity P Cygni profile on day 8.4 with velocity <1000 km s$^{-1}$, indicated by the blue absorption minimum.</figcaption></figure>

of He II $\lambda4686$ is very complicated with many lines blended together (see Figure 7), the flux in this region still seems to be too strong compared with the model spectra. This flux is therefore likely from the blueshifted emission of the fast-moving dense shell or the ejecta. The emission blueward of
<figure><img src="image_14.png" /><figcaption>**Figure 12.** Zoomed-in regions of $H\alpha$ (left panel), $H\beta$ (middle panel), and $H\gamma$ (right panel) profile from 5.5 to 14.5 days. A broad P Cygni profile ($v > 5000$ km s$^{-1}$) starts to emerge in $H\beta$ and $H\gamma$ (but not $H\alpha$) from day 8.4.</figcaption></figure>

He II λ4686 even strengthens in the first few days (see the right panel of Figure 7), supporting the fact that we are seeing progressively more He II λ4686 emission from the dense shell, which appears blueshifted at infinity due to optical-depth effects. The blueshifted emission eventually disappears after a week because He II recombines as the temperature drops.

During the CSM-dominated phase, in the higher-resolution spectrum at 1.58 days we are able to distinguish He II λ6560 from the blue wing of $H\alpha$, as shown in Figure 18. The ionized $H\alpha$ luminosity can be used to estimate the mass-loss rate before explosion (E. O. Ofek et al. 2013). To do so, we follow the steps presented by J. Zhang et al. (2023). We fit the $H\alpha$ profile with three components: a narrow Gaussian component for He II λ6560 (dashed cyan line), a narrow Gaussian component for $H\alpha$ (dashed green line), and a wide Lorentzian component for both $H\alpha$ and He II λ6560 (dashed blue line) as shown in Figure 18. We only adopt one wide component for both $H\alpha$ and He II

$\lambda6560$ because it is impossible to distinguish the two wide components given the small difference in central wavelength but very wide wings. We estimate the intrinsic FWHM of $H\alpha$ to be 80 ± 10 km s$^{-1}$, after removing by quadrature the instrumental FWHM of ~102 km s$^{-1}$ (estimated from the night-sky emission lines) from the original measurement of ~129 km s$^{-1}$. But note that the intrinsic FWHM of $H\alpha$ measured at this time may have already been changed because of radiative acceleration, so it is more appropriate to consider this estimate as an upper limit of the intrinsic FWHM. The flux ratio between narrow $H\alpha$ and He II λ6560 is ~6:4. The wide Lorentzian component gives an FWHM of ~2400 km s$^{-1}$, much broader than the narrow component. According to the relation $L_{H\alpha} \approx 2 \times 10^{39} \dot{M}_{0.01\text{v}_{\text{w,500}}\beta r_{15}}^2 \text{erg s}^{-1}$ (E. O. Ofek et al. 2013), we estimate the narrow $H\alpha$ luminosity to be $2.04 \times 10^{38} \text{erg s}^{-1}$. At 1.58 days we estimate the inner CSM distance (before it is swept up by the ejecta) to be
<figure><img src="image_15.png" /><figcaption>**Figure 13.** Zoomed-in regions of He I λ5876 (left) and He I λ7065 (right) from 5.5 to 14.5 days. The He I λ5876 profile show obvious absorption at a velocity of ~7000 km s⁻¹ around day 8.</figcaption></figure>

~1.4 × 10¹⁴ cm (see method below), so the mass-loss rate is found to be $\dot{M} \approx 4 \times 10^{-4} M_{\odot} \text{yr}^{-1}$. But note that the broad Hα component carries much more flux<fnref n="21" /> than the narrow Hα component; therefore, if counting the entire Hα luminosity, the estimated mass-loss rate would also be six times higher, $\dot{M} \approx 2.4 \times 10^{-3} M_{\odot} \text{yr}^{-1}$. This value is lower than the mass-loss rate adopted by the r1w6 model of $\dot{M} = 0.01 M_{\odot} \text{yr}^{-1}$, likely because the calculation ignored the optical-depth effects in Hα.

The Hα profile broadens quickly after the first 2 days, as shown in the middle panel of Figure 6. The FWHM measured at 2.4 days and 3.4 days was ~500 km s⁻¹. Interestingly, the FWHM measured from the wide Lorentzian component remains the same at ~2400 km s⁻¹, but is still much broader than the narrow component.

After the first week of the SN explosion, the Hα starts to show a P Cygni profile with a minimum absorption velocity of ~1000 km s⁻¹, as shown in Figure 11 (see also left panel of Figure 19). The P Cygni profile persists for more than a week (during this period, both the observed and r1w6 model spectra are mostly featureless as shown in Figure 8). The broadening of Hα from ~100 km s⁻¹ to ~1000 km s⁻¹ is unlikely to be caused by any outbursts or enhanced winds immediately before the explosion, but rather by radiative acceleration<fnref n="22" /> as

suggested by E. A. Zimmerman et al. (2024). However, the strong radiative acceleration mechanism requires the wind to have relatively low density, which means there is likely a big drop in the density profile following the shock breakout from the dense CSM. Since the P Cygni feature of Hα was detected as late as 14.5 days (after which the CSM would be swept up by the ejecta), by assuming an ejecta velocity of 10,000 km s⁻¹ (as inferred from the highest Hα velocity in Figure 15), one can calculate a distance of ~1.3 × 10¹⁵ cm; the low-density CSM surrounding the progenitor was present at least up to this distance. Similarly, one can also try to estimate the distance of the dense CSM from the time of shock breakout, which is roughly equal to the time of the temperature peak, because once the shock breaks out, the CSM starts to cool. From the curve of temperature evolution shown in Figure 9, we estimate the peak time to be day 2.5, thus giving a distance of ~2.2 × 10¹⁴ cm. This is about six times the progenitor radius from the model.

The drop in the CSM radial density profile indicates an enhanced episode of mass loss at a certain time prior to the final core collapse. Assuming a wind velocity of 80 km s⁻¹ (see above and Figure 18), the promptly ejected matter would require ~320 days to reach a distance of ~2.2 × 10¹⁴ cm. Since the typical RSG wind velocity is much smaller (~20 km s⁻¹) than 80 km s⁻¹, the actual delay time between the precursor event and the SN explosion could thus be longer, up to a few years. This could be consistent with the eruptive mass-loss scenario given by D. Hiramatsu et al. (2023), who favor eruptions releasing 0.3–1 $M_{\odot}$ of the envelope at ~1 yr before explosion, although other possibilities are also discussed by D. Hiramatsu et al. (2023). However, such a high-mass eruption before the explosion could be at odds with the appearance of SN IIn-like spectral signatures because the photons cannot escape (see L. Dessart & W. V. Jacobson-Galán 2023); thus, a gradually enhanced episode of mass loss is more plausible, though other possibilities also exist (N. Soker 2023; J. Fuller & D. Tsuna 2024). We also remark that the enhanced mass loss before the explosion inferred from the early-time spectral evolution of SN 2023ixf is corroborated by the semianalytic fits to its early chromatic luminosity evolution (G. Li et al. 2024). The latter demand the presence of a dense CSM shell confined within a distance of a few 10¹⁴–10¹⁵ cm from the progenitor star, which is most likely to be produced through elevated mass loss at a rate of ~10⁻³–10⁻² $M_{\odot} \text{yr}^{-1}$ a few years before the SN explosion.

At the same (CDS-dominated) phase, the He I lines (which disappeared earlier) also start showing broad P Cygni profiles around day 8, as shown in Figures 10 and 13. He I λ5876 has a velocity of 7000 km s⁻¹ (He I λ7065 is not very significant), similar to the velocity of Hβ, Hγ, and Hδ at the same phase. The emergence of broad He I is also predicted by the r1w6 model (Figure 10, red spectra) as the ejecta cool down and proceed to the recombination phase, though the model predicts the line to be present earlier (emerging on day 5 and disappearing by day 9).

On day 22, the broad P Cygni profile is well-developed and dominates the spectra. By this time, since the r1w6 model has no corresponding spectra at similar phases, we plot the shock-powered model Pwr1e41 (L. Dessart & D. J. Hillier 2022) instead, as shown in Figure 14. In the Pwr1e41 model, a dense shell was placed at 11,000 km s⁻¹ (which is very close to the ejecta velocity of ~10,000 km s⁻¹ measured from Hα) to

<footnote n="21">²¹ Measured to be about 10 times more, but only 60% is contributed by the broad Hα component, and 40% is contributed by the broad He II λ6560 component, assuming the flux ratio is the same as for the narrow component.</footnote>

<footnote n="22">²² An alternative explanation could be that we are seeing broad, boxy emission from the interaction with lower-density CSM, and that emission partially fills the Hα absorption. If the CSM is asymmetric and clumpy, this broad emission could be bumpy.</footnote>
<figure><img src="image_16.png" /><figcaption>**Figure 14.** Kast spectra of SN 2023ixf after three weeks (from days 22.5 to 120) in the ejecta-dominated phase, during which the spectra show many broad P Cygni features that are seen in typical normal SNe II. The red spectra were from the Pwr1e41 model at a nearby phase for comparison; see text for details.</figcaption></figure>

mimic the interaction of CSM at the earliest times and then a constant power of $10^{41}$ erg s$^{-1}$. Since the model was developed in 1D, both the dense shell and the shock power are distributed in a spherical shell at 11,000 km s$^{-1}$. We adopt the Pwr1e41 model here because it matches best with the observed spectra, including most of the major P Cygni features.

We also note that starting from day 84.5, a box-shaped component underlying the Hα profile emerges as indicated by the notch on its right side (also see A. Kumar et al. 2025). The boxy profile is an indication of the continued interaction between the ejecta and the CSM confined within a radially extended shell (A. Kumar et al. 2025). However, the edge of the boxy shape in the observed spectra shows a velocity of only ~7000 km s$^{-1}$ (see more details in Figure 19, right panel), much smaller than the ejecta velocity of ~10,000 km s$^{-1}$ measured from Hα, which suggests an asymmetric dense shell or shock power (asymmetric wind). The Pwr1e41 model also clearly predicts the boxy profile at a higher velocity of 11,000 km s$^{-1}$ (model input parameter), on both the red and blue sides. On the blue side, the box-shaped signature in the model spectra is a strong absorption feature at

11,000 km s$^{-1}$, which persists in the model spectra since day 15 (see Figure 19, right panel). However, such a strong feature is not clearly seen in the observed spectra, at least not before day 97 (there is a weak absorption feature after day 97 at similar velocity, but it is unclear whether it is related to the boxy profile). Instead, there is a boxy shape at ~6500 km s$^{-1}$, similar to the red side. If these two boxy profiles on both the blue and red sides are connected (since they are at similar velocity), it would suggest a possible axial symmetry, as in a bipolar explosion.<fnref n="23" /> In any case, since their velocities are much smaller than the ejecta velocity, it is an indication of an asymmetric dense shell, which may be related to the polarization. In fact, polarization of SN 2023ixf has been detected and reported at early phases (S. S. Vasylyev et al. 2023, 2025; A. Singh et al. 2024; M. Shrestha et al. 2025), and this further indicates that the CSM is indeed asymmetric. The asymmetry may help explain the observed lower velocity, as

<footnote n="23">²³ Note that it could also be possible that the CDS of SN 2023ixf is actually only at ~7000 km s$^{-1}$, but this conflicts with the velocity measured from Hα of ~10,000 km s$^{-1}$.</footnote>
<figure><img src="image_17.png" /><figcaption>**Figure 15.** Velocity evolution measured from the absorption minimum of the broad P Cygni profile. Dashed black and solid black lines show the Fe II λ5169 velocity evolution from C. P. Gutiérrez et al. (2017, G17) and T. Faran et al. (2014, F14), respectively. The evolution curve after three weeks (power-law shape) is typical of normal SNe II. However, the velocities of Hβ, Hγ, and Hδ (emerging earlier than Hα) at early times actually rise, which is not commonly seen in normal SNe II.</figcaption></figure>

we might be seeing only the projected velocity along our line of sight, which is much lower than the model CDS velocity at 11,000 km s⁻¹.

The nebular spectra can help us estimate the progenitor mass by comparing the observed spectra with the model spectra. We obtained several nebular spectra of SN 2023ixf at different phases. The model spectra were constructed based on the ejecta/progenitor models given by L. Dessart et al. (2021) and T. Sukhbold et al. (2016). Figure 20 shows the comparison at four epochs<fnref n="24" />: 207.7 days, 264.9 days, 299.9 days, and 362.6 days. It appears that the $15 M_\odot$ model matches best with the observed nebular spectrum (followed by the $12 M_\odot$ model) for almost all the epochs. This is also consistent with the direct progenitor analysis by S. D. Van Dyk et al. (2024), who constrained the initial mass of the progenitor candidate to be from $12 M_\odot$ to $15 M_\odot$. Comparing the four phases, we notice that the spectral flux decreased for all the lines except [Ca II] λλ7291, 7323 (middle of bottom zoomed-in panels for each phase in Figure 20), which actually became stronger relative to Ca II λλλ8498, 8542, 8662 (right side of bottom zoomed-in panels for each phase in Figure 20). This is reasonable; as the ejecta grow more optically thin, the forbidden lines of Ca II strengthen relative to the permitted Ca II triplet. Also, note that the Fe II emission below 5500 Å is much stronger than in the model spectra; this might come from supersolar metallicity, or greater ⁵⁶Ni mass, or (more likely) extra emission from interaction. Similarly, the nebular [Ca II] λλ7291, 7323 doublet on days 264.9 and 299.9 is also stronger than in the model spectra, which generally implies greater ⁵⁶Ni since the line forms primarily from the ⁵⁶Ni-rich material.

Most lines also show broad emission, implying a sizable CDS contributing Hα and perhaps some Ca II. In addition, the “bridge” between Hα and [O I] λλ6300, 6364 is much stronger in the nebular spectra (left side of bottom zoomed-in panels for each phase in Figure 20), a clear signature of a broad Hα emission from interaction. This means that some Hα emission is present at high blueshifted velocities, while it seems less present on the opposite, redshifted/receding part of the ejecta, linking to evidence of asymmetry at earlier times (i.e., asymmetry still prevails far from the star, although it is unclear whether this is the same geometry). In any case, the nebular spectrum until 362.6 days and even 441.5 days suggests that SN 2023ixf continues to interact (ever since shock breakout) with CSM produced by wind mass loss prior to the explosion.

## 5. Distance Measurement

M101 is a unique galaxy for testing our current methods to measure extragalactic distances. We have precise distance measurements from Cepheids (W. L. Freedman et al. 2001; A. Saha et al. 2006; B. J. Shappee & K. Z. Stanek 2011; A. G. Riess et al. 2022), TRGB (L. Rizzi et al. 2007; B. J. Shappee & K. Z. Stanek 2011; R. L. Beaton et al. 2019; G. S. Anand et al. 2022), and SNe Ia (J. Vinkó et al. 2012; A. G. Riess et al. 2022). The most recent measurements from those three techniques are all in good agreement. Using Cepheids, A. G. Riess et al. (2022) obtained a distance modulus of 29.194 ± 0.03 mag; with TRGB, G. S. Anand et al. (2022) obtained 29.08 ± 0.05 mag; and with SN 2011fe, A. G. Riess et al. (2022) measured 29.04 ± 0.12 mag.

However, the distance to M101 can also be derived using SN 2023ixf. Even if SNe II display a large range of peak luminosities, it has been demonstrated that they are excellent distance indicators. Using theoretical and empirical methods to calibrate them (T. de Jaeger & L. Galbany 2023), we can measure distances with a precision of ~15% (T. de Jaeger et al. 2020).

We focus our effort on the standard candle method (SCM; M. Hamuy et al. 2001), which is the most common and most accurate empirical technique used to derive SN II distances. Developed by M. Hamuy et al. (2001), it is based on the correlation between the SN absolute magnitude and two observables: the expansion velocity and the color. The expansion velocity is generally determined from the blueshift of Fe II λ5169 or Hβ P Cygni features. Using this technique, we obtain a distance of 28.67 ± 0.14 mag. This value is smaller than those derived from other techniques (Cepheids, TRGB, SNe Ia). The difference can be explained by the presence of strong CSM interaction (D. Hiramatsu et al. 2023), which boosts the luminosity above that of a normal SN II. The SCM is not able to correct for this effect, so the derived M101 distance from SN 2023ixf is likely too small. However, even if the distance is somewhat inaccurate, we can try to include SN 2023ixf as a calibrator to derive the Hubble–Lemaître constant through the distance-ladder technique. Adding SN 2023ixf to the 13 calibrators from T. de Jaeger et al. (2022), and using its Cepheid-based distance, we obtain $H_0 = 73.1_{-3.5}^{+3.8} \text{ km s}^{-1} \text{ Mpc}^{-1}$ (statistical only) instead of $H_0 = 75.4_{-3.7}^{+3.8} \text{ km s}^{-1} \text{ Mpc}^{-1}$ (T. de Jaeger et al. 2022). As expected, the uncertainties decrease because the number of calibrators increases (13 versus 14), and the new value is

<footnote n="24">²⁴ Note that L. Dessart et al. (2021) only showed model spectra at 350 days; model spectra at other epochs were computed as part of this work.</footnote>
<figure><img src="image_18.png" /><figcaption>**Figure 16.** Kast spectra of SN 2023ixf in the nebular phase when the ejecta start to become optically thin; strong forbidden emission lines start to emerge, including [O I], Mg I], and Ca II. Spectra shown in red are from the Pwr1e41 model at a nearby phase for comparison.</figcaption></figure>
<figure><img src="image_19.png" /><figcaption>**Figure 17.** Montage showing the evolution of the ejecta/CSM properties from the radiation hydrodynamical simulation of model r1w6 developed by D17, together with the regions of formation of various lines as computed with CMFGEN. The shaded area represents regions with $\tau_{es} > 2/3$.</figcaption></figure>

<figure><img src="image_20.png" /><figcaption>**Figure 18.** H$\alpha$ profile from our higher-resolution spectra at day 1.58, where He II $\lambda$6560 is distinguished from H$\alpha$. Three components are used to fit the profile: a narrow Gaussian component for He II $\lambda$6560 (dashed cyan line), a narrow Gaussian component for H$\alpha$ (dashed green line), and a wide Lorentzian component for both H$\alpha$ and He II $\lambda$6560 (dashed blue line).</figcaption></figure>

<figure><img src="image_21.png" /><figcaption>**Figure 19.** Left: early-time evolution (first two weeks) of the H$\alpha$ profile compared with the r1w6 model (L. Dessart et al. 2017) at similar phases. The FWHM of H$\alpha$ increases quickly after first two days. After a week, H$\alpha$ starts to show a P Cygni profile with a minimum absorption velocity of ~1000 km s$^{-1}$, which lasts for more than a week. Right: later-time evolution (after three weeks) of H$\alpha$ compared with the Pw1e41 model (L. Dessart & D. J. Hillier 2022) at similar phases. The H$\alpha$ profile starts to show a boxy shape at day 84.5, clearly on the red side, but with an edge velocity of only ~7000 km s$^{-1}$, much smaller than the ~10,000 km s$^{-1}$ ejecta velocity. This is direct evidence for an asymmetric dense shell.</figcaption></figure>

smaller because SN 2023ixf is more luminous than the average SN II.

More recently, a theoretical method called the tailored expanding photosphere method has been developed by C. Vogl (2020). It is based on the expanding photosphere method (R. P. Kirshner & J. Kwan 1974) but avoids the systematic distance uncertainties from the dilution factors. This method has been applied to several objects, yielding distance uncertainties smaller than 5% (G. Csörnyei et al. 2023). However, it cannot be used for SN 2023ixf because the radiative-transfer codes of C. Vogl (2020) do not include the strong CSM interaction seen in this object.

## 6. Conclusion and Discussion

We present both photometric and spectroscopic observations of SN 2023ixf covering from early times to the nebular phase. The light curves show that SN 2023ixf has a relatively short plateau phase (~72 days), and is either in the fast-decliner (IIL) subclass or a transitional event between SNe IIP and IIL, indicating an H-rich envelope of lower mass in the progenitor before explosion. It reached a peak V-band absolute magnitude of $-18.2 \pm 0.07$ mag, thus putting it at the bright end of SNe II.

Optical spectra show that SN 2023ixf transitioned from Type IIn to a typical Type II SN after three weeks. Early-time spectra of SN 2023ixf exhibit strong, narrow emission lines
<figure><img src="image_22.png" /><figcaption>**Figure 20.** Spectra of SN 2023ixf in the nebular phase at days 207.7, 264.9, 299.9, and 362.6, compared with models (L. Dessart et al. 2021) having different masses. Each bottom panel shows zoomed-in regions of the three strongest lines to reveal details of the comparisons.</figcaption></figure>
from the ionized CSM. We identified many species that produce these lines and found that most of them are from H and He, including the He II series from energy states of $n = 4$, 5, and 6. The emission features weakened after the first week and the spectra appear to be blue and quasi-featureless in the second week, and can be fitted with a blackbody. After ~3 weeks, the spectra are similar to those of other SNe II, with strong P Cygni features, and the SN entered the nebular phase after about six months.

We compare observed spectra of SN 2023ixf with various model spectra in order to understand the physics behind SN 2023ixf. There is likely to be dense CSM surrounding the progenitor up to a distance of ~2.2 × 10¹⁴ cm, followed by lower-density CSM to distances of at least ~1.3 × 10¹⁵ cm. Shortly before exploding, the progenitor may have had a mass-loss rate of up to $\dot{M} \approx 2.4 \times 10^{-3} M_{\odot} \text{yr}^{-1}$. Our nebular spectra match best with a 15 $M_{\odot}$ model spectrum.

SN 2023ixf is used as a distance indicator by fitting the light curves; we obtain a distance modulus of 28.67 ± 0.14 mag, slightly smaller than that derived from other techniques (Cepheids, TRGB, SN Ia), possibly because of strong CSM interaction in SN 2023ixf. By including SN 2023ixf as a calibrator and using its Cepheid distance, we obtain $H_0 = 73.1_{-3.50}^{+3.68} \text{ km s}^{-1} \text{Mpc}^{-1}$ (statistical only), slightly smaller (because SN 2023ixf is more luminous than the average SN II) but still in good agreement with $H_0 = 75.4_{-3.7}^{+3.8} \text{ km s}^{-1} \text{Mpc}^{-1}$ given by T. de Jaeger et al. (2022).

## Acknowledgments

A major upgrade of the Kast spectrograph on the Shane 3 m telescope at Lick Observatory, led by Brad Holden, was made possible through gifts from the Heising-Simons Foundation, William and Marina Kast, and the University of California Observatories. KAIT and its ongoing operation were made possible by donations from Sun Microsystems, Inc., the Hewlett-Packard Company, AutoScope Corporation, Lick Observatory, the U.S. National Science Foundation, the University of California, the Sylvia & Jim Katzman Foundation, and the TABASGO Foundation. Research at Lick Observatory is partially supported by a gift from Google. Some of the data presented herein were obtained at the W. M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California, and NASA; the observatory was made possible by the generous financial support of the W. M. Keck Foundation. We appreciate the expert assistance of the staffs at Lick and Keck Observatories.

We thank the following Lick Observatory Graduate Workshop participants for help obtaining Nickel photometry on two nights and Shane spectra on one night: Fangyi Cao, Casey Carlile, Madalyn Johnson, Zhexing Li, Ye Lin, Patrick Maloney, Michael McDonald, Pradyumna Sadhu, Loraine Sandoval Ascencio, Sogul Sanjaripour, Niloofar Sharei, Hurum Maksora Tohfa, and Michael Wozniak.

Generous financial support was provided to A.V.F.'s supernova group at U.C. Berkeley by the Christopher R. Redlich Fund, Steven Nelson, Sunil Nagaraj, Landon Noll, Sandra Otellini, Gary and Cynthia Bengier, Clark and Sharon Winslow, Alan Eustace and Kathy Kwan, William Draper, Timothy and Melissa Draper, Briggs and Kathleen Wood, and Sanford Robertson (S.S.V. is a Steven Nelson Graduate Fellow in Astronomy, K.C.P. was a Nagaraj-Noll-Otellini Graduate

Fellow in Astronomy, W.Z. is a Bengier-Winslow-Eustace Specialist in Astronomy, T.G.B. is a Draper-Wood-Robertson Specialist in Astronomy, Y.Y. was a Bengier-Winslow-Robertson Fellow in Astronomy). Numerous other donors to his group and/or research at Lick Observatory include Michael Antin, Charlie Baxter and Jinee Tao, Duncan Beardsley, Jim Connelly and Anne Mackenzie, Curtis and Shelley Covey, Byron and Allison Deeter, Paul and Silvia Edwards, Arthur and Cindy Folker, Heidi Gerster, Harvey Glasser, Charles and Gretchen Gooding, Thomas and Dana Grogan, Judith and Timothy Hachman, Alan and Gladys Hoefer, George Hume, Stephen and Catherine Imbler, Michael Kast and Rebecca Lyon, Lata Krishnan and Ajay Shah, Walter and Karen Loewenstern, Rand Morimoto and Anna Henderson, Edward Oates, Doug and Emilie Ogden, Jonathan and Susan Reiter, Paul Robinson, Catherine Rondeau, Laura Sawczuk and Luke Ellis, Stanley Schiffman, Richard and Betsey Sesler, Hans Spiller, Justin and Seana Stephens, Ilya Strebulaev and Anna Dvornikova, David Turner III and Joanne Turner, Diane Tokugawa and Alan Gould, David and Joanne Turner, David and Malin Walrod, Gerald and Virginia Weiss, Janet Westin and Mike McCaw, Byron and Nancy Wood, David and Angie Yancey, and others. Y.Y.'s research is partially supported by the Tsinghua University Dushi Program. A.S. acknowledges support from NASA/HST grant GO-17216. L.D. acknowledges access to the HPC resources of TGCC under the allocation 2023—A0150410554 on Irene-Rome made by GENCI, France.

**Software:** Astropy (Astropy Collaboration et al. 2013, 2018, 2022), IDL Astronomy user's library (W. B. Landsman 1993).

## ORCID iDs

WeiKang Zheng https://orcid.org/0000-0002-2636-6508

Luc Dessart https://orcid.org/0000-0003-0599-8407

Alexei V. Filippenko https://orcid.org/0000-0003-3460-0103

Yi Yang (杨轶) https://orcid.org/0000-0002-6535-8500

Thomas G. Brink https://orcid.org/0000-0001-5955-2502

Thomas de Jaeger https://orcid.org/0000-0001-6069-1139

Sergiy S. Vasylyev https://orcid.org/0000-0002-4951-8762

Schuyler D. Van Dyk https://orcid.org/0000-0001-9038-9950

Kishore C. Patra https://orcid.org/0000-0002-1092-6806

Wynn V. Jacobson-Galán https://orcid.org/0000-0002-3934-2644

Gabrielle E. Stewart https://orcid.org/0009-0000-2503-140X

Efrain Alvarado, III https://orcid.org/0009-0005-8159-8490

Adam J. Burgasser https://orcid.org/0000-0002-6523-9536

Evan M. Carrasco https://orcid.org/0009-0005-5223-1606

Elma Chuang https://orcid.org/0000-0001-9984-5131

Asia DeGraw https://orcid.org/0009-0001-2794-8278

Elinor L. Gates https://orcid.org/0000-0002-3739-0423

Cooper Jacobus https://orcid.org/0000-0002-0498-2019

Preethi R. Karpoor https://orcid.org/0000-0002-1480-9041

Neil Pichay https://orcid.org/0009-0009-7665-6827

Jon Rees https://orcid.org/0000-0002-5376-3883

R. Michael Rich https://orcid.org/0000-0003-0427-8387

Nathan R. Sandford https://orcid.org/0000-0002-7393-3595
Alessandro Savino https://orcid.org/0000-0002-1445-4877

Emma Softich https://orcid.org/0000-0002-1420-1837

Christopher A. Theissen https://orcid.org/0000-0002-9807-5435

Edgar P. Vidal https://orcid.org/0009-0002-2209-4813

William Wu https://orcid.org/0009-0000-0753-4345

References

Anand, G. S., Tully, R. B., Rizzi, L., Riess, A. G., & Yuan, W. 2022, ApJ, 932, 15

Anderson, J. P., González-Gaitán, S., Hamuy, M., et al. 2014, ApJ, 786, 67

Astropy Collaboration, Price-Whelan, A. M., Lim, P. L., et al. 2022, ApJ, 935, 167

Astropy Collaboration, Price-Whelan, A. M., Sipőcz, B. M., et al. 2018, AJ, 156, 123

Astropy Collaboration, Robitaille, T. P., Tollerud, E. J., et al. 2013, A&A, 558, A33

Beaton, R. L., Seibert, M., Hatt, D., et al. 2019, ApJ, 885, 141

Bonanos, A. Z., & Boumis, P. 2016, A&A, 585, A19

Bostroem, K. A., Pearson, J., Shrestha, M., et al. 2023, ApJL, 956, L5

Bruch, R. J., Gal-Yam, A., Yaron, O., et al. 2023, ApJ, 952, 119

Chufarin, V., Potapov, N., Ionov, I., et al. 2023, TNSAN, 150, 1

Csörnyei, G., Vogl, C., Taubenberger, S., et al. 2023, A&A, 672, A129

de Jaeger, T., & Galbany, L. 2023, arXiv:2305.17243

de Jaeger, T., Galbany, L., Riess, A. G., et al. 2022, MNRAS, 514, 4620

de Jaeger, T., Zheng, W., Stahl, B. E., et al. 2019, MNRAS, 490, 2799

Dessart, L. 2024, arXiv:2405.04259

Dessart, L., Gutiérrez, C. P., Ercolino, A., Jin, H., & Langer, N. 2024, A&A, 685, A169

Dessart, L., & Hillier, D. J. 2005, A&A, 439, 671

Dessart, L., & Hillier, D. J. 2010, MNRAS, 405, 2141

Dessart, L., & Hillier, D. J. 2022, A&A, 660, L9

Dessart, L., Hillier, D. J., & Audit, E. 2017, A&A, 605, A83

Dessart, L., Hillier, D. J., Audit, E., Livne, E., & Waldman, R. 2016, MNRAS, 458, 2094

Dessart, L., Hillier, D. J., Sukhbold, T., Woosley, S. E., & Janka, H. T. 2021, A&A, 652, A64

Dessart, L., & Jacobson-Galán, W. V. 2023, A&A, 677, A105

Dickinson, D., Milisavljevic, D., Garretson, B., et al. 2025, ApJ, 984, 71

Faran, T., Poznanski, D., Filippenko, A. V., et al. 2014, MNRAS, 442, 844

Filippenko, A. V. 1982, PASP, 94, 715

Filippenko, A. V. 1989, AJ, 97, 726

Filippenko, A. V. 1991a, in Supernovae, ed. S. E. Woosley (New York: Springer-Verlag), 467

Filippenko, A. V. 1991b, in European Southern Observatory Conf. and Workshop Proc. 37, ed. I. J. Danziger & K. Kjaer (Garching: ESO)

Filippenko, A. V. 1997, ARA&A, 35, 309

Filippenko, A. V., Li, W. D., Treffers, R. R., & Modjaz, M. 2001, in ASP Conf. Ser. 246, IAU Colloq. 183: Small Telescope Astronomy on Global Scales, ed. B. Paczynski, W.-P. Chen, & C. Lemme (San Francisco, CA: ASP), 121

Filippenko, A. V., Zheng, W., & Yang, Y. 2023, TNSAN, 123, 1

Freedman, W. L., Madore, B. F., Gibson, B. K., et al. 2001, ApJ, 553, 47

Fuller, J., & Tsuna, D. 2024, OJAp, 7, 47

Fulton, M. D., Nicholl, M., Smith, K. W., et al. 2023, TNSAN, 124, 1

Gal-Yam, A., Arcavi, I., Ofek, E. O., et al. 2014, Natur, 509, 471

Galbany, L., Hamuy, M., Phillips, M. M., et al. 2016, AJ, 151, 33

Gutiérrez, C. P., Anderson, J. P., Hamuy, M., et al. 2017, ApJ, 850, 89

Hamann, N. 2023, TNSAN, 127, 1

Hamuy, M., & Pinto, P. A. 2002, ApJL, 566, L63

Hamuy, M., Pinto, P. A., Maza, J., et al. 2001, ApJ, 558, 615

Henden, A., Krajci, T., & Munari, U. 2012, IBVS, 6024, 1

Hillier, D. J., & Dessart, L. 2012, MNRAS, 424, 252

Hillier, D. J., & Dessart, L. 2019, A&A, 631, A8

Hiramatsu, D., Howell, D. A., Moriya, T. J., et al. 2021, ApJ, 913, 55

Hiramatsu, D., Tsuna, D., Berger, E. 2023, ApJL, 955, L8

Hosseinzadeh, G., Farah, J., Shrestha, M., et al. 2023, ApJL, 953, L16

Itagaki, K. 2023, TNSTR, 1, 2023

Jacobson-Galán, W. V., Dessart, L., Davis, K. W., et al. 2024, ApJ, 970, 189

Jacobson-Galán, W. V., Dessart, L., Jones, D. O., et al. 2022, ApJ, 924, 15

Jacobson-Galán, W. V., Dessart, L., Margutti, R., et al. 2023, ApJL, 954, L42

Jencson, J. E., Pearson, J., Beasor, E. R., et al. 2023, ApJL, 952, L30

Kilpatrick, C. D., Foley, R. J., Jacobson-Galán, W. V., et al. 2023, ApJL, 952, L23

Kirshner, R. P., & Kwan, J. 1974, ApJ, 193, 27

Koltenbah, B. 2023, TNSAN, 144, 1

Kumar, A., Dastidar, R., Maund, J. R., Singleton, A. J., & Sun, N.-C. 2025, MNRAS, 538, 659

Landolt, A. U. 1992, AJ, 104, 340

Landsman, W. B. 1993, in ASP Conf. Ser. 52, Astronomical Data Analysis Software and Systems II, ed. R. J. Hanisch, R. J. V. Brissenden, & J. Barnes (San Francisco, CA: ASP), 246

Leonard, D. C., Filippenko, A. V., Barth, A. J., & Matheson, T. 2000, ApJ, 536, 239

Li, G., Hu, M., Li, W., et al. 2024, Natur, 627, 754

Li, W., Filippenko, A. V., Chornock, R., & Jha, S. 2003, PASP, 115, 844

Limeburner, S. 2023, TNSAN, 128, 1

Mao, Y., Zhang, M., Cai, G., et al. 2023, TNSTNS, 130, 1

Miller, J., & Stone, R. 1994, The Kast Double Spectrograph 66, Univ. California https://mthamilton.ucolick.org/techdocs/instruments/kast/Tech%20Report%2066%20KAST%20Miller%20Stone.pdf

Morozova, V., Piro, A. L., Renzo, M., et al. 2015, ApJ, 814, 63

Niu, Z., Sun, N.-C., Maund, J. R., et al. 2023, ApJL, 955, L15

Ofek, E. O., Lin, L., Kouveliotou, C., et al. 2013, ApJ, 768, 47

Perley, D. A., Gal-Yam, A., Irani, I., & Zimmerman, E. 2023, TNSAN, 119, 1

Perley, D. A., & Irani, I. 2023, TNSAN, 120, 1

Pledger, J. L., & Shara, M. M. 2023, ApJL, 953, L14

Qin, Y.-J., Zhang, K., Bloom, J., et al. 2024, MNRAS, 534, 271

Ransome, C. L., Habergham-Mawson, S. M., Darnley, M. J., et al. 2021, MNRAS, 506, 4715

Riess, A. G., Yuan, W., Macri, L. M., et al. 2022, ApJL, 934, L7

Rizzi, L., Tully, R. B., Makarov, D., et al. 2007, ApJ, 661, 815

Saha, A., Thim, F., Tammann, G. A., Reindl, B., & Sandage, A. 2006, ApJS, 165, 108

Sanders, N. E., Soderberg, A. M., Gezari, S., et al. 2015, ApJ, 799, 208

Schlaflay, E. F., & Finkbeiner, D. P. 2011, ApJ, 737, 103

Schlegel, E. M. 1990, MNRAS, 269, 244

Sgro, L. A., Esposito, T. M., Blaclard, G., et al. 2023, RNAAS, 7, 141

Shappee, B. J., & Stanek, K. Z. 2011, ApJ, 733, 124

Shrestha, M., DeSoto, S., Sand, D. J., et al. 2025, ApJL, 982, L32

Silverman, J. M., Foley, R. J., Filippenko, A. V., et al. 2012, MNRAS, 425, 1789

Singh, A., Teja, R. S., Moriya, T. J., et al. 2024, ApJ, 975, 132

Smartt, S. J. 2009, ARA&A, 47, 63

Smartt, S. J. 2015, PASA, 32, e016

Smith, N., Pearson, J., Sand, D. J., et al. 2023, ApJ, 956, 46

Soker, N. 2023, RAA, 23, 081002

Soraisam, M. D., Szalai, T., Van Dyk, S. D., et al. 2023, ApJ, 957, 64

Stahl, B. E., Zheng, W., de Jaeger, T., et al. 2019, MNRAS, 490, 3882

Stathakis, R. A., & Sadler, E. M. 1991, MNRAS, 250, 786

Stetson, P. B. 1987, PASP, 99, 191

Sukhbold, T., Ertl, T., Woosley, S. E., Brown, J. M., & Janka, H. T. 2016, ApJ, 821, 38

Tartaglia, L., Sand, D. J., Groh, J. H., et al. 2021, ApJ, 907, 52

Teja, R. S., Singh, A., Basu, J., et al. 2023, ApJL, 954, L12

Terreran, G., Jacobson-Galán, W. V., Groh, J. H., et al. 2022, ApJ, 926, 20

Terreran, G., Jerkstrand, A., Benetti, S., et al. 2016, MNRAS, 462, 137

Tody, D. 1986, Proc. SPIE, 627, 733

Valenti, S., Howell, D. A., Stritzinger, M. D., et al. 2016, MNRAS, 459, 3939

Van Dyk, S. D. 2017, RSPTA, 375, 20160277

Van Dyk, S. D., Li, W., & Filippenko, A. V. 2003, PASP, 115, 1289

Van Dyk, S. D., Peng, C. Y., King, J. Y., et al. 2000, PASP, 112, 1532

Van Dyk, S. D., Srinivasan, S., Andrews, J. E., et al. 2024, ApJ, 968, 27

Van Dyk, S. D., Zheng, W., Maund, J. R., et al. 2019, ApJ, 875, 136

Vasylyev, S. S., Dessart, L., Yang, Y., et al. 2025, arXiv:2505.03975

Vasylyev, S. S., Yang, Y., Filippenko, A. V., et al. 2023, ApJL, 955, L37

Vinkó, K., Sárneczky, K., Takáts, K., et al. 2012, A&A, 546, A12

Vogl, C. 2020, PhD thesis, Technical Univ. Munich, Germany

Xiang, D., Mo, J., Wang, L., et al. 2024, SCPMA, 67, 219514

Yamanaka, M., Fujii, M., & Nagayama, T. 2023, PASJ, 75, L27

Yaron, O., Bruch, R., Chen, P., et al. 2023, TNSAN, 133, 1

Yaron, O., Perley, D. A., Gal-Yam, A., et al. 2017, NatPh, 13, 510

Zhang, J., Lin, H., Wang, X., et al. 2023, SciBu, 68, 2548

Zhang, K., Kennedy, D., Oostermeyer, B., Bloom, J., & Perley, D. A. 2023, TNSAN, 125, 1

Zimmerman, E. A., Irani, I., Chen, P., et al. 2024, Natur, 627, 759