100+ Nebular SNe Ia

Clearing the Smoke: Nebular Spectra of 100+ Type Ia Supernovae Exclude Single Degenerate Progenitors

M. A. Tucker, B. J. Shappee, P. J. Vallely, K. Z. Stanek, J. Prieto, J. Botyanszki, C. S. Kochanek, J. P. Anderson, J. Brown, L. Galbany, T. W.-S. Holoien, E. Y. Hsiao, S. Kumar, H. Kuncarayakti, N. Morrell, M. M. Phillips, M. D. Stritzinger, and Todd A. Thompson
Institute for Astronomy, University of Hawai‘i at Manoa, 2680 Woodlawn Dr., Honolulu, Hi 96822
Department of Astronomy, The Ohio State University, 140 West 18th Avenue, Columbus, OH 43210, USA
Center for Cosmology and AstroParticle Physics (CCAPP), The Ohio State University, 191 W. Woodruff Avenue, Columbus, OH 43210, USA
Núcleo de Astronomía de la Facultad de Ingeniería y Ciencias, Universidad Diego Portales, Av. Ejército 441, Santiago, Chile
Millennium Institute of Astrophysics, Santiago, Chile
Physics Department, University of California, Berkeley, CA 94720, USA
European Southern Observatory, Alonso de Cordova 3107 Casilla 19001, Vitacura, Santiago, Chile
PITT PACC, Department of Physics and Astronomy, University of Pittsburgh, Pittsburgh, PA 15260, USA
Carnegie Observatories, 813 Santa Barbara Street, Pasadena, CA 91101, USA
Department of Physics, Florida State University, 77 Chieftan Way, Tallahassee, FL 32306, USA
Finnish Centre for Astronomy with ESO (FINCA), FI-20014 University of Turku, Finland
Tuorla Observatory, Department of Physics and Astronomy, FI-20014 University of Turku, Finland
Las Campanas Observatory, Carnegie Observatories, Casilla 601, La Serena, Chile
Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark
E-mail: tuckerma@hawaii.eduDOE CSGF Fellow
Accepted XXX. Received YYY; in original form ZZZ
Abstract

We place statistical constraints on Type Ia supernova (SN Ia) progenitors using 226 nebular phase spectra of SNe Ia. We find no evidence of stripped companion emission in any of the nebular phase spectra. Upper limits are placed on the amount of mass that could go undetected in each spectrum using recent hydrodynamic simulations. With these null detections, we place an observational upper limit on the fraction of SNe Ia that are produced through the classical H-rich non-degenerate companion scenario of . Additionally, we set a tentative upper limit on He star progenitor scenarios of , although further theoretical modelling is required. As part of our analysis, we also derive a Nebular Phase Phillips Relation, which approximates the brightness of a SN Ia in the nebular phase using the peak magnitude and decline rate parameter .

keywords:
supernovae – general; galaxies – distances and redshifts
pubyear: 2019pagerange: Clearing the Smoke: Nebular Spectra of 100+ Type Ia Supernovae Exclude Single Degenerate ProgenitorsB.3

1 Introduction

Type Ia supernovae (SNe Ia) are utilised across many astronomical disciplines, including the extragalactic distance scale, dark energy studies, and Galactic chemical evolution. Despite their prevalence, the origins of SNe Ia are still unclear even after decades of study. The general consensus is that they are explosions of carbon/oxygen (C/O) white dwarfs (Hoyle & Fowler, 1960) with fairly homogenous properties. For example, the magnitude of SNe Ia at peak is well constrained (, e.g.; Folatelli et al., 2010a), and, after correcting for light curve decline and color, they have an intrinsic scatter of mag (e.g., Fig. 19, Folatelli et al., 2010a). Many formation mechanisms for SNe Ia have been proposed to reproduce this level of unformity, which can be grouped into two main categories: the double degenerate (DD) and single degenerate (SD) scenarios (see Maoz et al., 2014; Livio & Mazzali, 2018, for reviews on SNe Ia progenitors).

The DD scenario consists of two degenerate stars, usually C/O white dwarfs, which induce a SNe Ia through accretion, collision, or merger. This can occur due to gravitational wave emission (Tutukov & Yungelson, 1979; Iben & Tutukov, 1984; Webbink, 1984), collision/violent merger due to perturbations by external bodies (Thompson, 2011; Katz & Dong, 2012; Shappee & Thompson, 2013; Pejcha et al., 2013; Antognini et al., 2014), accretion from a low-mass white dwarf onto a smaller, higher-mass white dwarf (Taam, 1980; Livne, 1990; Pakmor et al., 2012), or a "double detonation" where an accreted helium layer detonates and drives the core to detonate (Woosley & Weaver, 1994; Fink et al., 2010; Kromer et al., 2010). Due to the intrinsic faintness of both components in these systems, observational confirmation of DD systems is exceptionally difficult (e.g., Rebassa-Mansergas et al., 2018). Some progress has been made on this front, such as bimodal emission in the nebular phase (Dong et al., 2015a; Vallely et al., 2019) and possible hyper-velocity remnants (Shen et al., 2018). However, most of the evidence for DD systems comes from the exclusion of SD progenitors (e.g., Shappee et al., 2017).

The SD scenario involves a WD with a nearby non-degenerate companion (Whelan & Iben, 1973; Nomoto, 1982; Yoon & Langer, 2003), usually undergoing Roche Lobe overflow (RLOF). The WD accumulates material until reaching critical mass and then explodes. This critical mass is typically considered the Chandrasekhar mass (), although sub- explosions, including double detonation scenarios, are also possible (e.g., Livne & Arnett, 1995). There are several predicted observational signatures of the SD degenerate scenario due to the interaction of the ejecta/explosion and the donor star (Wheeler et al., 1975), including effects on the rising SN Ia light curve (Kasen, 2010), soft X-ray emission in the accretion phase (Lanz et al., 2005; Woods et al., 2018), surviving companions with anomalous characteristics (e.g., Canal et al., 2001; Shappee et al., 2013b), and the amount of Ni decay products synthesized in the explosion (e.g., Röpke et al., 2012; Shappee et al., 2017).

One of the most promising signatures of a RLOF companion to an exploding WD are emission lines produced by material stripped/ablated from the non-degenerate companion (e.g., Wheeler et al., 1975; Marietta et al., 2000; Mattila et al., 2005; Pan et al., 2012), observable in nebular-phase spectra once the SN Ia has faded considerably and become optically thin. For example, Boehner et al. (2017) simulated stripping from red giant (RG), main sequence (MS), and sub-giant (SG) stars, finding approximately , , and , respectively, of stripped mass. Botyánszki et al. (2018) converted these estimates into expected H luminosities and found that the emitted H luminosity does not vary linearly with amount of stripped companion mass, which had been the assumption of previous studies (e.g. Leonard, 2007; Shappee et al., 2013a), but instead the relation is closer to exponential. Additionally, the H emission is powered by the SN Ia and roughly follows the bolometric luminosity.

In this work we compile a comprehensive sample of SNe Ia nebular spectra spanning after explosion to search for the expected emission from stripped/ablated material. We find no such emission in any spectrum in our sample, and place new or updated stripped/ablated mass constraints for each SN Ia. The entirety of similar work in the literature totals 25 SNe Ia (Mattila et al., 2005; Leonard, 2007; Shappee et al., 2013a; Lundqvist et al., 2013, 2015; Maguire et al., 2016; Graham et al., 2017; Shappee et al., 2018; Sand et al., 2018a; Holmbo et al., 2018; Dimitriadis et al., 2019; Tucker et al., 2018), a fraction of the sample analyzed in this work. All SNe Ia included in this study are listed in Table 6.

We outline our data sources and reduction techniques, including absolute flux calibration, in §2. In §3, we discuss our methodology in searching for and placing limits on material stripped from a RLOF companion. Our upper limits on stripped material are provided in §4, and our findings are discussed in the context of SNe Ia formation in §5. Included in §5 are discussions about peculiar SNe Ia and their role in our study, plus the curious case of ASASSN-18tb which exhibits broad H emission in its semi-nebular spectrum (Kollmeier et al., 2019).

2 Data Sources and Reduction

Our sample of 226 spectra of SNe Ia comes from the 38 instruments on 27 telescopes listed in Table 1. All spectroscopically peculiar SNe Ia are included except for those exhibiting signatures of circumstellar material (SNe Ia-CSM). These SNe Ia exhibit H emission, but the velocity and magnitude of the emission is inconsistent with material stripped from a nearby companion. Instead, these SN Ia appear to have exploded in a dense circumstellar environment (e.g., SN 2002ic, Wang et al., 2004). We impose the following criteria when selecting SNe Ia nebular spectra:

  • Obtained between 200 and 500 days after explosion to maintain consistency with the models of Botyánszki et al. (2018), assuming a typical rise time of (Firth et al., 2015).

  • Cover of at least one H or He line in Table 2.

  • Have at least one method of absolute flux calibration, outlined in §2.3.

The complete list of new and archival spectra is provided in Table 8. Additionally, we include new and archival photometry to supplement our spectral data and analysis. Early phase photometry ( after maximum light) is used in deriving the photometric properties of each SN Ia using the photometric fitting code SNooPy (Burns et al., 2011), including time of maximum (), the decline rate parameter , extinction along line of sight, and the distance modulus. Late- and nebular-phase photometry are used for flux calibrating the nebular spectra and deriving a Nebular Phase Phillips Relation (NPPR). The NPPR approximates the nebular magnitude of a SN Ia given its peak magnitude and decline rate, calibrated to an extensive sample of new and archival SNe Ia photometry. A complete description of the NPPR, its derivation and usage is provided in Appendix A.

Telescope Abbrev. Instrument Abbrev. Ref.
Australian National University 2.3m ANU2.3m Wide-Field Spectrograph WiFeS Dopita et al. (2007, 2010) 7
Calar Alto 2.2m CA2.2m Calar Alto Faint Object Spectrograph CAFOS 1
Calar Alto 3.5m CA3.5m Multi-Object Spectrograph at Calar Alto MOSCA 2
Danish 1.54m D1.54m Danish Faint Object Spectrograph and Camera DFOSC Andersen et al. (1995) 1
du Pont Telescope duPont Wide Field Reimaging CCD Camera WFCCD 4
Boller and Chivens Spectrograph BC 1
ESO 1.5m ESO1.5m Boller and Chivens Spectrograph BC 2
ESO 3.6m ESO3.6m ESO Faint Object Spectrograph and Camera EFOSC1/2 Buzzoni et al. (1984) 11
Himalayan Chandra Telescope HCT Himalayan Faint Object Spectrograph HFOSC 2
Hubble Space Telescope HST Faint Object Spectrograph FOS 1
Isaac Newton Telescope INT Faint Object Spectrograph (1st Gen.) FOS1 Breare et al. (1987) 2
Gemini North/South GN/S Gemini Multi-Object Spectrograph GMOS Hook et al. (2004) 12
Gran Telescopio Canarias GTC Optical System for Imaging and low-Intermediate-Resolution Integrated Spectroscopy OSIRIS Cepa (2010) 1
Keck I KeckI Low Resolution Imaging Spectrograph LRIS Oke et al. (1995) 27
Keck II KeckII DEep Imaging Multi-Object Spectrograph DEIMOS Faber et al. (2003) 9
Echelette Imager and Spectrograph ESI Sheinis et al. (2002) 3
Large Binocular Telescope LBT Multi-Object Double Spectrograph MODS Pogge et al. (2010) 10
Magellan Baade Telescope Baade Inamori-Magellan Areal Camera and Spectrograph IMACS Dressler et al. (2011) 3
Magellan Echellette Spectrograph MagE Marshall et al. (2008) 2
Magellan Clay Telescope Clay Low Dispersion Survey Spectrograph LDSS 5
Multiple Mirror Telescope MMT Blue Channel Spectrograph BCS Angel et al. (1979) 6
New Technology Telescope NTT ESO Multi-Mode Instrument EMMI D’Odorico (1990) 1
SOFI Moorwood et al. (1998) 1
Palomar 200-inch P200 Double Spectrograph DBSP Oke & Gunn (1982) 4
Shane 3m Telescope Shane3m Kast Spectrograph KAST Silverman et al. (2013) 11
Southern African Large Telescope SALT Robert Stobie Spectrograph RSS Buckley et al. (2006) 3
Subaru Sub OH-Airglow Suppressor/Cooled Infrared Spectrograph and Camera for OHS CISCO Motohara et al. (2002) 3
Faint Object Spectrograph and Camera FOCAS Kashikawa et al. (2002) 2
Tillinghast 1.5m Till FAst Spectrograph for the Tillinghast telescope FAST Fabricant et al. (1998) 7
Telescopio Nazionale Galileo TNG Device Optimized for LOw RESolution DOLORES Molinari et al. (1999) 2
Very Large Telescope VLT FOcal Reducer and low dispersion Spectrograph FORS1/2 Appenzeller et al. (1998) 43
Multi-Unit Spectroscopic Explorer MUSE Bacon et al. (2010) 8
XSHOOTER XSH Vernet et al. (2011) 16
William Herschel Telescope WHT Intermediate dispersion Spectrograph and Imaging System ISIS Jorden (1990) 5
ACAM Benn et al. (2008) 1
Faint Object Spectrograph (2nd Gen.) FOS2 Breare et al. (1987) 6
Unknown Unknown 1
Total 27 38 226
Table 1: All telescopes and instruments utilised in this work. If a reference could not be found for a given instrument, the corresponding instrument website is provided in the table notes.

2.1 New Spectra and Photometry

We present 14 new nebular-phase spectra of 13 SNe Ia, of which 10 have no prior published nebular spectra. These spectra were acquired in our ongoing study of SNe Ia progenitors, taken with MagE on Baade, MUSE on the VLT, and WFCCD on duPont (see Table 1 for telescope and instrument designations). For the new spectra presented here, each spectrum was reduced using telescope and instrument-specific pipelines, if available, otherwise typical IRAF tasks were used. The spectra acquired with MagE/Baade were reduced with a pipeline provided by the Carnegie Observatories111http://code.obs.carnegiescience.edu/mage-pipeline (Kelson et al., 2000; Kelson, 2003), with the exception of standard star calibrations and stitching together each echellette spectrum, which was done with custom Python routines. For newly presented MUSE data acquired as part of the AMUSING survey (Galbany et al., 2016), spectra were extracted in a 1" circular aperture at the SN Ia location using the PyMUSE package (Pessa et al., 2018), and corrected for host galaxy contributions using a background annulus extending from 2" to 3".

For absolute flux calibrations, we also include nebular photometry for any SNe Ia in our sample. This includes new observations and reproccessed archival images for which we could not find a published magnitude. New photometry includes -band images taken with FORS2, -band images from MODS1, and images from WFCCD. Archival imaging includes imaging from FORS1/2 and imaging from EFOSC2 (Table 14). All images are bias subtracted and flat-field corrected before performing aperture photometry with the IRAF apphot task. For targets with , photometry from the Pan-STARRS Stack Object catalog222http://archive.stsci.edu/panstarrs/stackobject/search.php (Chambers et al., 2016; Flewelling et al., 2016) was used in calibrating the images, otherwise Gaia DR2 photometry (Gaia Collaboration et al., 2016, 2018; Riello et al., 2018) was used. When transforming reported magnitudes to other photometric systems, Tonry et al. (2012) and Evans et al. (2018) were used for Pan-STARRS and Gaia, respectively. The only exceptions to this procedure are the -band FORS2/VLT images, which are calibrated using the reported photometric zeropoints333https://www.eso.org/observing/dfo/quality/FORS2/qc/zeropoints/zeropoints.html.

2.2 Archival Spectra and Photometry

The primary sources of our archival spectra and photometry are the Berkeley SuperNova Ia Program444http://heracles.astro.berkeley.edu/sndb/ (BSNIP, Silverman et al., 2012, 2013), the Center for Astrophysics (CfA) Supernova Data Archive555https://www.cfa.harvard.edu/supernova/SNarchive.html (Riess et al., 1999; Jha et al., 2006; Matheson et al., 2008; Blondin et al., 2012), the Carnegie Supernova Project666http://csp.obs.carnegiescience.edu/ (CSP, Hamuy et al., 2006; Folatelli et al., 2010a; Contreras et al., 2010; Stritzinger et al., 2011; Folatelli et al., 2013; Krisciunas et al., 2017; Phillips et al., 2019), the 100IAs project (Dong et al., 2018a), the ANU WiFeS SuperNovA Program (AWSNAP; Childress et al., 2016), and the All-Sky Automated Survey for SuperNovae (ASAS-SN, Shappee et al., 2014a; Holoien et al., 2017a, b, c, 2019, Vallely et al., Chen et al., in prep). The majority of the publicly available data were retrieved using the Open Supernova Catalog (OSC, Guillochon et al., 2017) and the Weizmann Interactive Supernova data REPository (WISeREP, Yaron & Gal-Yam, 2012). All data provided by these sources are already reduced with the exception of precise spectral flux calibration, which we outline in §2.3. Additionally, we supplement these sources with archival data obtained from telescope databases, including the Keck Observatory Archive777https://koa.ipac.caltech.edu/ (KOA), the ESO Science Archive Facility 888http://archive.eso.org/cms.html (ESO SAF), the Isaac Newton Group Archive999http://casu.ast.cam.ac.uk/casuadc/ingarch/query and the Gemini Observatory Archive101010https://archive.gemini.edu/ (GOA). Information on all the spectra in this study is presented in Table 8.

Data reduction and calibration was performed as uniformly as possible across all sources of spectra. Data retrieved from public archives were already reduced, with the exception of absolute flux calibration. The reduction of data retrieved from telescope archives was generally less complete. All spectra retrieved from the ESO SAF were already reduced (excluding flux corrections) with the exception of FORS1/2 data. For any ESO SAF data reduction, both spectroscopy and photometry, we used the ESO SAF esorex data reduction pipeline (Freudling et al., 2013).

Spectra obtained from the KOA and GOA were not reduced prior to retrieval and had to be manually reduced. Recent LRIS spectra were reduced using Lpipe111111http://www.astro.caltech.edu/~dperley/programs/lpipe.html, while older LRIS and DEIMOS data were reduced using the LowRedux/XIDL pipeline121212http://www.ucolick.org/~xavier/LowRedux/. Gemini North/South GMOS spectra were reduced with the GMOS Data Reduction Cookbook131313http://ast.noao.edu/sites/default/files/GMOS_Cookbook/.

We manually reduced any unreduced spectra for which no pipeline exists using standard IRAF141414http://iraf.noao.edu/ procedures. Images were flat-fielded and bias-subtracted using archival calibration images taken near the epoch of observation, and wavelength calibrated with arc lamp exposures. Spectrophotometric standard star observations were used to correct for telescope/instrumental artefacts, atmospheric effects, and to place each spectrum on a reliable relative flux scale.

2.3 Accurate Flux Calibration

For our analysis in §3, the spectra must be on a reliable absolute flux scale. While calibrating spectra with spectrophotometric standard stars places these spectra on a dependable relative flux scale, slit losses, atmospheric conditions, and other effects can cause the resulting spectra to deviate from an absolute flux scale. To scale a spectrum to the absolute scale, we employed Eq. 7 from Fukugita et al. (1996) to calculate synthetic photometry from the spectra. The spectra are then scaled so that the synthetic photometry matches the observed photometry. There were several different sources of photometry used to calibrate the spectra. In order of preference and reliability, with accuracy estimates in given parentheses:

  1. For spectra with acquisition images taken at the time of observation, we scale the entire spectrum to match these photometric observations, usually in the or filters ().

  2. If acquisition images are unavailable, we next tried to use photometry within of the spectral observations. Photometry in all available filters within this temporal limit were used in the flux calibration ().

  3. If no photometric data was available within , we searched for photometry within . If at at least 3 photometric data points fell within this time span, we linearly interpolated to estimate the magnitude at the time of the spectral observation ().

  4. If none of these were available, the nebular magnitude was estimated with the NPPR and used to calibrate the spectrum (see Appendix A, ).

We required of the filter’s transmission curve be covered by the observed spectrum for viable calibrations. If only a single filter was available, the entire spectrum was scaled to match the observation. If two filters were available for flux calibration, a simple linear fit was applied to the scale factors. If filters were available, we use spline fits with fixed endpoints to ensure a robust flux correction across the entire spectrum. After placing the spectrum on an absolute flux scale, we correct for host galaxy and Milky Way reddening using the derived from the light curve fits. We implement a Fitzpatrick (1999) extinction law and a Schlegel et al. (1998) Milky Way dust map for our reddening corrections. We assume unless stated otherwise (see Appendix B).

3 Searching for Emission from a Stripped Companion

Prior to the work of Botyánszki et al. (2018), the majority of unbound mass limits in the literature utilised the work of Mattila et al. (2005) and Leonard (2007) to compute stripped mass limits from comparing observed spectra to expected H luminosities. Several subsequent studies have adopted these methodologies in their work (e.g., Maguire et al., 2016; Graham et al., 2017) with notable success in ruling out hydrogen-rich companions. Yet the models of Mattila et al. (2005) had several shortcomings in observational implementation. In particular, Leonard (2007) assumed a linear scaling between the amount of unbound companion mass and the corresponding H luminosity.

Botyánszki et al. (2018), using the MS38 model (a main sequence star undergoing RLOF) from Boehner et al. (2017), instead found the emitted H luminosity scales exponentially with the amount of stripped mass. Additionally, Botyánszki et al. (2018) computed a simplified helium-star model, where all the stripped mass from the MS38 model is replaced with helium instead of Solar abundance material. This is not a true helium star model, as helium star companions are expected to have lower amounts of stripped mass than their hydrogen-rich counterparts and a modestly different velocity distribution (e.g. Pan et al., 2012), but it provides a starting point for calculating limits on the amount of unbound helium in a SN Ia spectrum.

While the models of Botyánszki et al. (2018) clarify the mass-luminosity scaling issue and expand to helium emission, they share two other shortcomings with the models of Mattila et al. (2005): the requirement of using H to constrain the amount of unbound mass, and only calculating the expected H luminosity at a single epoch (200 days post-explosion for Botyánszki et al., 2018 and 350 days post-maximum for Mattila et al., 2005). In the following subsections we discuss our stripped mass limits given these limitations.

3.1 Expanding on these Models

Line [ erg/s] FWHM [Å]
H-rich Model
H 4341Å 0.271 14.5
H 4831Å 4.38 16.1
HeI-a 5875Å 4.27 19.6
H 6563Å 68.0 21.9
HeI-b 6678Å 2.24 22.3
HeI-c 1.08 10.5 36.0
Pa 1.281 14.6 42.7
Pa 1.875 14.6 62.5
HeI-d 2.06 8.48 68.7
He-rich Model
HeI-a 5875Å 8.26 19.6
HeI-b 6678Å 6.90 22.3
HeI-c 1.08 18.2 36.0
HeI-d 2.06 12.9 68.7
Table 2: Line luminosities for both the hydrogen-rich (H-rich) model and the helium-rich (He-rich) model corresponding to the MS38 and helium models from Botyánszki et al. (2018). Helium lines are given letter designations to ease identification in Table 12. FWHM refers to the expected FWHM of a line profile broadened by .

For SNe Ia with star-forming host galaxies, the region around H can be contaminated by narrow host galaxy H and NII emission lines which complicates setting limits on H emission. However, the unbound material has emission lines besides H, including H, H, and the Paschen series. Assuming roughly Solar metallicity, the stripped material will also exhibit prominent HeI lines in the optical and NIR (Botyánszki et al., 2018). We provide the luminosities for each of these lines in Table 2 at days for the hydrogen-rich (H-rich) model using the same MS38 model as Botyánszki et al. (2018). Additionally, we supply similar data for the simplified helium star model from Botyánszki et al. (2018), which we refer to as the He-rich model.

Figure 1: Peak luminosity of the FeIII line (red points) versus from the time-dependent SN Ia spectral models of Botyánszki & Kasen (2017) and the exponential fit (black line).
Figure 2: Nebular phase spectrum (black), continuum fit (red), and derived flux limits (blue) for the Baade/MagE +295 d spectrum of SN 2015F. The bottom panels show the regions near each possible emission line from Table 2 and correspond to the coloured boxes in the top panel. Gray shaded areas indicate masked spectral regions due to host galaxy contamination or instrumental effects, and vertical dashed lines indicate SNe Ia emission lines used to estimate the smoothing width for each spectrum (see §3). Similar spectral cutouts for all SNe Ia are included as supplementary figures (see Appendix B).

Botyánszki et al. (2018) estimated the line luminosity at 200 days as a function of the amount of stripped mass (). Table 2 provides the expected luminosity of various lines for . The dependence on the amount of stripped mass is well approximated by , where . Botyánszki et al. (2018) do not provide the time dependence of the line emission specifically, but note that the H emission is proportional to the FeIII emission over the  day period they consider. Utilising the synthetic spectra models from Botyánszki & Kasen (2017), we find the FeIII emission is well fit by an exponential (Fig. 1), which leads to an estimate for the line luminosity of

(1)

provided . This should hold well for the Balmer lines, and is at least a better approximation for the Paschen and HeI lines than assuming their luminosities are temporally constant.

3.2 Placing Statistical Limits on Stripped Mass

Once each spectrum is flux calibrated and corrected for the reddening, we place statistical limits on the presence of emission lines listed in Table 2, roughly following the methods of Leonard (2007). Each spectrum is rebinned to the approximate spectral resolution, and the spectral continuum is fit with a 2-order Savitsky-Golay polynomial (Press et al., 1992), excluding pixels within of line centers to prevent biasing our continuum fit, as done in previous studies (e.g., Maguire et al., 2016). However, since we are inspecting a multitude of lines for emission signatures, we apply our continuum model in velocity space instead of wavelength space to incorporate this modification.

No single continuum width adequately fits the continuum for all SNe Ia in our sample, especially considering the spectroscopic and temporal diversity. We tailored the continuum fit width for each spectrum based on the observed SN Ia expansion velocity, measured from the prominent emission lines in the spectrum. Since most of the major emission lines in nebular SNe Ia are blended to some extent (e.g., Mazzali et al., 2015, Fig. 5), we compute the weighted average from the fitted line profiles assuming a Gaussian emission profile + linear continuum. The lines considered for deriving the expansion velocity are the major FeII, FeIII, and CoIII lines indicated by the vertical dashed lines in Fig. 2. If the SNR of the spectrum is too low for the widths of at least 2 lines to be measured confidently, we assume a typical width of . For velocities lower than this value, we risk biasing our continuum fit to include possible weak emission, and implement as a strict lower bound. Additionally, since SNe Iax are known to have narrow line profiles in the nebular phase compared to typical SNe Ia (Foley et al., 2016), we adopt this lower bound for SNe-Iax as well. Because these velocities are simply a proxy for the width of the continuum fit, this method neglects the intricacies of SNe Ia emission profiles, especially since spectroscopically bi-modal SNe Ia are not uncommon (Dong et al., 2015a; Vallely et al., 2019). However, these complications are unimportant for our analysis, and we consider these simple velocity approximations adequate.

When applying the continuum fit to each spectrum, we minimize biasing our continuum by using clipping to exclude narrow host galaxy lines, telluric absorption, or instrumental artefacts. After fitting the continuum model to the data, we subtract off this continuum and inspect the residuals for emission line signatures from unbound companion material. For each line in Table 2, we compute bounds on the integrated line flux in each region similar to Eq. 4 from Leonard & Filippenko (2001). However, for flux calibrated spectra,

(2)

where is the upper limit on the integrated flux, is the corresponding upper limit on the equivalent width, is the continuum flux at wavelength , and is the RMS scatter around a normalised continuum. Eq. 2 can be re-written as

(3)

where is the RMS scatter of the spectrum around the continuum in flux units (erg s cm Å) and is the correction term for masked pixels (see §3.3). Our statistical limit may seem overly conservative but it does correspond to a line profile that would be visibly obvious (e.g., Fig. 3). Additionally, other studies have run injection-recovery tests to determine the true detection threshold for emission lines in SNe Ia nebula spectra and a purely statistical is difficult to recover (e.g., Sand et al., 2018a).

is then converted into a luminosity via the distance moduli. Distance moduli computed from the SN Ia light curves are used except where more reliable methods are available, such as Cepheid or Tip of the Red Giant Branch (TRGB) distances. Eq. 3 is inverted to numerically calculate a limit on , which we consider a conservative upper bound on the amount of mass removed from a non-degenerate companion undergoing RLOF. This is done for each H/He line, retaining the best mass limit for both the H-rich and He-rich models. Note that the strictest mass limit for each model can come from different spectra, as each spectrum will have varying amounts of contamination from host galaxy and telluric lines.

3.3 Mitigating Host Galaxy Emission and Other Contaminants

Due to the comprehensive nature of our sample, some spectra have poor quality, significant host-galaxy emission and/or other contaminants. Pixels affected by host galaxy emission, telluric absorption, or instrumental artefacts are masked in the ensuing flux limit calculation, ensuring only informative pixels are used in placing our flux upper limit. Masking these pixels also reduces the effective number of pixels used in the non-detection limit calculation and weakens our statistical limit. In Eq. 3 we include the masked pixel correction term from Tucker et al. (2018) to correct our limit to a more robust estimate. Concisely, the correction term is the fraction of unmasked line flux to total line flux (). Thus, masked pixels decrease and increase , but the effect is weighted by the location of the masked pixels relative to the line centre. For example, the masked narrow host galaxy H and H in the bottom panels of Fig. 2 have larger effects on than the masked [SII] line at Å since [SII] is on the outskirts of the HeI-a line profile. Masking is only implemented when the derived is not representative of the true flux limit due to contaminated pixels, we leave weak or minor contamination unmasked as it only solidifies our conservative flux limit and does not introduce extra steps in our analysis.

Another difficulty occurs when the expected emission line is blended with the edge of a steep SN spectral feature. This is especially problematic for 91bg-like and Iax SNe which have intrinsically narrow emission line profiles. If the continuum near H/He varies by more than the amplitude of our flux limit over its FWHM, we increase our flux limit to match the continuum level variation. This results in an unambiguous line profile that would be definitively detected and prevents questionable limits from being included in our statistical analysis.

Some spectra in our study have resolutions of order , which approaches the lower end of the expected stripped mass velocity distribution (e.g., Boehner et al., 2017). If broad, unresolved H emission was present in a spectrum, we confirm the host galaxy source with other typical galaxy emission lines such as [OII] (Å), [OIII] (, Å), [NII] (6548, 6583Å), and [SII] (6713, 6731Å). Any unresolved H emission with velocity widths had at least one other unresolved galaxy emission line in the spectrum, indicating the observed H emission was not from stripped material. Additionally, the recent discovery of broad H emission in ASASSN-18tb (Kollmeier et al., 2019) affirms our treatment of galaxy emission lines, as none of the galaxy emission lines discussed previously were present in the discovery spectrum (see §5).

4 Results

Figure 3: Randomly selected cutouts around H for a portion of the SNe Ia in this work, including the observed spectrum (black), the continuum fit (red), and the empirical line limit (purple). The scale for each spectrum is denoted in the top-left of each panel. Light grey areas mark masked regions (see text) and completely grey boxes signify SNe Ia with no spectra covering the wavelength range. Thick blue axes indicate this spectrum was used for the best limit provided in Table 12. Cutouts for all SNe Ia and all H/He lines are provided as supplementary material (see Appendix B).
Figure 4: Distribution of mass limits on stripped H-rich material for all SNe Ia in our sample. Colour-shaded areas indicate expected amounts of unbound mass for sub-giant (SG, blue), main sequence (MS, green), and red giant (RG, red) companions, taken from Marietta et al. (2000), Pan et al. (2012), and Boehner et al. (2017).
Figure 5: Similar to Fig. 4, except for the He-rich model. The magenta shaded region corresponds to stripped mass estimates from Liu et al. (2013a) and the cyan line marks the estimate from Pan et al. (2012).
Figure 6: Similar to Fig. 4, except using the models of Mattila et al. (2005). The assumed linear scaling between luminosity and stripped mass leads to higher derived .

We find no evidence of emission from stripped/ablated companion material in any of our nebular phase spectra. Fig. 2 provides an example Baade/MagE spectrum of SN 2015F at +295 days after maximum light, including the observed spectrum, the continuum fit, and the flux limits for each line. We provide a random selection of H flux limit cutouts in Fig. 3 and the spectral cutouts for all H and He lines are provided as supplementary material.

The distribution of stripped mass limits are shown for the H-rich and the He-rich cases in Figs. 4 and 5, respectively, with colour-shaded regions indicating the expected amounts of stripped mass from various studies in the literature. Fig. 6 shows the H-rich results using the methods and models of Mattila et al. (2005) and Leonard (2007) for comparison with previous estimates. Table 12 gives the phases, flux limits, and derived H-rich and He-rich mass limits for each SN Ia in our study.

We include the range of mass estimates from an H-rich RLOF companion in Figs. 4 and 6 as shaded regions for main-sequence (MS, blue), sub-giant (SG, green), and red giant (RG, red) companions taken from Marietta et al. (2000), Pan et al. (2012), and Boehner et al. (2017). We take as the minimum amount of mass stripped from a companion in the SD scenario, SNe Ia with are considered unlikely to have an H-rich SD progenitor system.

For the He-rich SD channel, only Pan et al. (2012) and Liu et al. (2013a) have published models. We include their expected values for mass stripped from a RLOF helium-star companion in Fig. 5 as the magenta shaded area (Liu et al., 2012) and the cyan line (Pan et al., 2012). However, there are several caveats when considering the He-rich model. The expected line luminosities given in Table 2 are for of stripped He-rich material, more mass than expected for a true He-donor star. We compare our mass limits to the dedicated He-rich models from Pan et al. (2012), Liu et al. (2013a) and take a limit of as our upper limit for He-rich SD systems.

Figure 7: Distribution of in this sample compared to the photometric sample from Ganeshalingam et al. (2010). As expected, the nebular sample is biased towards brighter and broader SNe Ia.
H-rich He-rich
() ()
Sample
Normal
91T-like
91bg-like
SC
Iax
Normal+91T
Normal+91bg
Normal+91T+91bg
Normal+91T+91bg+SC
All
Table 3: Statistics for each sample considered in our study (see §4). is the number of SNe Ia with and is the fractional upper limit on their occurrence. refers to the total number of SNe Ia in that sample.

If we assume SNe Ia with and exclude H-rich and He-rich SD progenitor systems, respectively, we can constrain the observed fraction of SD systems. Based on the non-detections in our sample, we can place observed upper limits on the fraction of SD SNe Ia. For a binomial distribution with trials and no successes, the upper limit at a confidence level can be expressed as

(4)

with the results for our sample provided in Table 3. and correspond to the and fractional upper limits on SD SNe Ia. For our null detections of unbound mass emission, we place statistical constraints on the fraction of SNe Ia that can form through the classical SD scenario for H-rich and He-rich companions. We do not consider SNe Ia with inadequate limits on “successes”, as the spectra do not show any evidence of the expected emission signatures, so these objects are simply omitted from our statistical analysis.

5 Discussion

With our updated modelling and comprehensive sample, we place strict constraints on the fraction SNe Ia that can form through the classical SD scenario. At most, of SNe Ia can stem from the H-rich formation channel, placing the majority of the production of SNe Ia on the DD channel, unless a modification on the SD scenario can prevent nearly all SNe Ia from exhibiting these expected H and He emission signatures such as the spin-up/spin-down scenario (Di Stefano et al., 2011). Considering the simplest case of only spectroscopically normal SNe Ia, we place a () upper limit on SD progenitors of (). The full statistical results are provided in Table 3, and we use the Normal+91T+91bg+SC sample as the most representative sample from our survey. Unfortunately, our sample prevents an analysis of under- versus over-luminous SNe Ia, as we are biased towards brighter SNe Ia (Fig. 7). This highlights the importance of volume-limited surveys such as 100IAs (Dong et al., 2015a). Still, these stringent constraints on the observed rate of SD SNe Ia provide strong evidence for the DD channel producing the majority of SNe Ia.

Some SNe Ia spectroscopic sub-classes, such as 91T- and 91bg-like, are thought to stem from the same basic mechanism as normal SNe Ia. Because these SNe Ia are thought to be on the edges of typical SN Ia formation, we compare the derived stripped mass limits to the same expected stripped mass values as normal SNe Ia. Our sample has 91bg-like and 91T-like SNe Ia, for which we place () upper limits on H-rich SD progenitors at () and (), respectively.

For SNe Iax and “Super Chandrasekhar” (SC) SNe Ia, it is worth discussing their characteristics and applicability. The Iax sub-type (Foley et al., 2013) is thought to stem from an entirely different formation mechanism and appear to never enter a nebular phase but instead have photospheric properties (Foley et al., 2016). Our study includes such systems: SNe 2002cx, 2005hk, 2008A, and 2012Z. Liu et al. (2013b) investigated the expected values of unbound mass for these systems if in a SD system, finding significantly lower values of compared to the typical range. All Iax SNe Ia in our sample have , so the statistics are unchanged if the more stringent mass limit is employed. However, even if material is unbound from non-degenerate donor stars in these SNe Ia, it is still unclear if this material would be visible at late times. For these reasons, our main statistical analysis excludes these objects.

Our sample also includes “Super Chandrasekhar" (SC) SNe Ia explosions (SNe 2006gz, 2007if, 2009dc, and SNF 20080723-012), where the inferred ejecta mass, , is higher than the Chandrasekhar mass of (e.g., Howell et al., 2006; Scalzo et al., 2019). The main SD channel formation theory of SC SNe Ia involves a WD above the Chandrasekhar mass spinning rapidly enough to sustain itself from collapse and explosion. The high angular momentum was gained during accretion from a non-degenerate companion onto the WD, which then slows over time and eventually explodes. In some scenarios, the spin-down of the WD is longer than the lifetime of its companion, leading to a SD SN Ia with no stripped mass emission. This mechanism has also been used to explain regular, non-SC SNe Ia (e.g. the spin-up/spin-down scenario, Di Stefano et al., 2011). However, conclusive evidence that the WD always outlives the donor star has not been presented, so we include these SNe Ia in our preferred sample.

For completeness and comparison to the literature, we also derive mass limits using the prior models of Marietta et al. (2000) and Mattila et al. (2005) which are shown in Fig. 6. Considering the same preferred Normal+91T+91bg+SC sample, we still rule out H-rich non-degenerate companions () for SNe Ia, corresponding to a () fractional upper limit of (). This result differs slightly from the upper limit provided in Table 3 due to the assumed linear (instead of exponential) scaling between stripped mass and emitted luminosity (e.g., Leonard, 2007). Regardless, the majority of SNe Ia cannot stem from the SD scenario unless there is a systematic flaw or unincorporated physical process shrouding or suppressing the expected emission.

In addition to the observational limitations discussed in §3, the companion-interaction models are developed for normal SNe Ia. The amount of mass stripped from a RLOF companion for over- and under-luminous SNe Ia may differ from the results of Boehner et al. (2017), although the magnitude of these effects is currently unexplored. Specific stripped mass models exist for Iax SNe (Liu et al., 2013b), but do not exist for the other sub-types included in this paper. Furthermore, quantifying these effects is beyond the scope of this work.

There are other studies which place quantitative or qualitative limits on the fraction of SD progenitor systems using a range of wavelengths and techniques (e.g., Gilfanov & Bogdán, 2010; Hayden et al., 2010; Bianco et al., 2011; Brown et al., 2012; Chomiuk et al., 2016, this work). Most of these studies focus on WD+RG systems, as these are the easiest to observationally detect. Each study individually does not definitively rule out SD SNe Ia progenitors, however, when considered as a whole it is clear that most SNe Ia cannot form through the classical SD scenario. Thus, the DD scenario must account for the majority of normal SNe Ia. However, detecting and characterising double WD binaries is exceptionally difficult (e.g., Rebassa-Mansergas et al., 2018). Even accounting for observational difficulties, it is still unclear if the current formulation of the DD channel, including extensions, can provide of observed SNe Ia.

At present, there is one exception to the lack of nebular H emission, ASASSN-18tb (Kollmeier et al., 2019). The origin of the emission is unclear, as the observed flux is lower than typical SNe Ia-CSM but the inferred stripped mass value of is inconsistent with the SD model. However, these models have limitations (see §3) and the current simulations do not investigate the effects of subluminous explosions, which likely produce lower amounts of stripped mass. Late-onset SNe Ia-CSM have also been discovered (Dilday et al., 2012; Graham et al., 2019) and the velocity width of H emission increases with time in these systems, indicating this could also be similar to CSM objects although a SD progenitor system is definitely possible. Regardless, the true origin of the H observed in ASASSN-18tb is unknown and the early spectrum () lies outside the temporal bounds of our study (). Thus, we cannot include this object in our statistical analysis, although if this unique object is not a member of the SNe Ia-CSM class this discovery highlights the inherent rarity of such SNe Ia.

In summary, we find no evidence of emission from stripped/ablated companion material in any of nebular phase spectra. With this null result, we calculate the upper limit on the fraction of SNe Ia that can form through the SD channel for both H-rich and He-rich companions. At confidence, we find H-rich companions are restricted to and He-rich companions restricted to . These results provide strict constraints on the SD channel of SNe Ia formation, requiring of SNe Ia to stem from the DD channel.

Facilities: duPont, Magellan, Very Large Telescope

Software: Python2.7, astropy (The Astropy Collaboration et al., 2018), astroquery (Ginsburg et al., 2019) numpy, scipy, PyMUSE (Pessa et al., 2018), SpectRes (Carnall, 2017), extinction151515https://github.com/kbarbary/extinction, SExtractor (Bertin & Arnouts, 1996), Montage161616http://montage.ipac.caltech.edu/, Lpipe, IDL8.6, LowRedux, IRAF, SNooPy (Burns et al., 2011)

Acknowledgements

We thank K. Maguire, K. Graham, K. Motohara, K. Maeda, S. Taubenberger, T. Diamond, G. Dimitriadis, C. Ashall, and D. Sand for supplying nebular spectra. We thank A. Laity for assisting with the KOA data search and retrieval. Additionally, we thank C. Auge, G. Anand, A. Payne, and O. Graur for useful conversations about the project. We thank P. Chen from providing several SNe Ia light curves prior to publication.

M.A.T. acknowledges support from the United States Department of Energy through the Computational Sciences Graduate Fellowship (DOE CSGF). M.D.S. is supported in part by a generous grant (13261) from VILLUM FONDEN and a project grant from the Independent Research Fund Denmark. Support for J.L.P. is provided in part by FONDECYT through the grant 1191038 and by the Ministry of Economy, Development, and Tourism’s Millennium Science Initiative through grant IC120009, awarded to The Millennium Institute of Astrophysics, MAS.

This work is Based on (in part) observations collected at the European Organisation for Astronomical Research in the Southern Hemisphere under ESO programmes 0100.D-0191(A), 0101.D-0173(A), 0102.D-0287(A), and 096.D-0296(A). This paper includes data gathered with the 6.5 meter Magellan Telescopes located at Las Campanas Observatory, Chile. The CSP has been supported by the National Science Foundation under grants AST0306969, AST0607438, AST1008343, AST1613426, and AST1613472.

This paper made use of the modsIDL spectral data reduction reduction pipeline developed in part with funds provided by NSF Grant AST-1108693.

This research made use of Montage. It is funded by the National Science Foundation under Grant Number ACI-1440620, and was previously funded by the National Aeronautics and Space Administration’s Earth Science Technology Office, Computation Technologies Project, under Cooperative Agreement Number NCC5-626 between NASA and the California Institute of Technology.

This work has made use of data from the European Space Agency (ESA) mission Gaia (https://www.cosmos.esa.int/gaia), processed by the Gaia Data Processing and Analysis Consortium (DPAC, https://www.cosmos.esa.int/web/gaia/dpac/consortium). Funding for the DPAC has been provided by national institutions, in particular the institutions participating in the Gaia Multilateral Agreement.

The Pan-STARRS1 Surveys (PS1) and the PS1 public science archive have been made possible through contributions by the Institute for Astronomy, the University of Hawaii, the Pan-STARRS Project Office, the Max-Planck Society and its participating institutes, the Max Planck Institute for Astronomy, Heidelberg and the Max Planck Institute for Extraterrestrial Physics, Garching, The Johns Hopkins University, Durham University, the University of Edinburgh, the Queen’s University Belfast, the Harvard-Smithsonian Center for Astrophysics, the Las Cumbres Observatory Global Telescope Network Incorporated, the National Central University of Taiwan, the Space Telescope Science Institute, the National Aeronautics and Space Administration under Grant No. NNX08AR22G issued through the Planetary Science Division of the NASA Science Mission Directorate, the National Science Foundation Grant No. AST-1238877, the University of Maryland, Eotvos Lorand University (ELTE), the Los Alamos National Laboratory, and the Gordon and Betty Moore Foundation.

Based on observations obtained at the Gemini Observatory acquired through the Gemini Observatory Archive and processed with the Gemini PyRAF package, which is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the National Research Council (Canada), CONICYT (Chile), Ministerio de Ciencia, Tecnología e Innovación Productiva (Argentina), and Ministério da Ciência, Tecnologia e Inovação (Brazil).

This paper makes use of data obtained from the Isaac Newton Group of Telescopes Archive which is maintained as part of the CASU Astronomical Data Centre at the Institute of Astronomy, Cambridge.

The LBT is an international collaboration among institutions in the United States, Italy and Germany. LBT Corporation partners are: The University of Arizona on behalf of the Arizona Board of Regents; Istituto Nazionale di Astrofisica, Italy; LBT Beteiligungsgesellschaft, Germany, representing the Max-Planck Society, The Leibniz Institute for Astrophysics Potsdam, and Heidelberg University; The Ohio State University, and The Research Corporation, on behalf of The University of Notre Dame, University of Minnesota and University of Virginia.

References

  • Abazajian et al. (2005) Abazajian K., et al., 2005, AJ, 129, 1755
  • Abolfathi et al. (2018) Abolfathi B., et al., 2018, The Astrophysical Journal Supplement Series, 235, 42
  • Adelman-McCarthy et al. (2008) Adelman-McCarthy J. K., et al., 2008, The Astrophysical Journal Supplement Series, 175, 297
  • Akerlof et al. (2007) Akerlof C., et al., 2007, Central Bureau Electronic Telegrams, 1059, 2
  • Aksenov (1981) Aksenov E. P., 1981, International Astronomical Union Circular, 3580, 1
  • Altavilla et al. (2004) Altavilla G., et al., 2004, MNRAS, 349, 1344
  • Amanullah et al. (2014) Amanullah R., et al., 2014, ApJ, 788, L21
  • Andersen et al. (1995) Andersen J., et al., 1995, The Messenger, 79, 12
  • Angel et al. (1979) Angel J. R. P., et al., 1979, in The MMT and the Future of Ground-Based Astronomy, Proceedings of a Symposium held to makr the dedication of the Multiple Mirror Telescope at the Mount Hopkins Observatory, Arizona on May 9, 1979. Edited by Trevor C. Weekes. SAO Special Report #385, 1979., p.87. p. 87
  • Antognini et al. (2014) Antognini J. M., et al., 2014, MNRAS, 439, 1079
  • Anupama et al. (2005) Anupama G. C., et al., 2005, A&A, 429, 667
  • Appenzeller et al. (1998) Appenzeller I., et al., 1998, The Messenger, 94, 1
  • Arbour et al. (1999) Arbour R., et al., 1999, International Astronomical Union Circular, 7156, 1
  • Arbour et al. (2004) Arbour R., et al., 2004, International Astronomical Union Circular, 8406, 1
  • Arcavi et al. (2014) Arcavi I., et al., 2014, The Astronomer’s Telegram, 6661, 1
  • Ardeberg & de Groot (1973) Ardeberg A., et al., 1973, A&A, 28, 295
  • Argyle et al. (1994) Argyle R. W., et al., 1994, International Astronomical Union Circular, 5976, 3
  • Armstrong & Schwartz (1999) Armstrong M., et al., 1999, International Astronomical Union Circular, 7108, 1
  • Ayani & Yamaoka (1998) Ayani K., et al., 1998, International Astronomical Union Circular, 6878, 2
  • Ayani et al. (1998) Ayani K., et al., 1998, International Astronomical Union Circular, 6905, 1
  • Bacon et al. (2010) Bacon R., et al., 2010, in Proceedings of the SPIE, Volume 7735, id. 773508 (2010).. , doi:10.1117/12.856027
  • Barbon et al. (1972) Barbon R., et al., 1972, International Astronomical Union Circular, 2411, 1
  • Barbon et al. (1982) Barbon R., et al., 1982, A&A, 116, 35
  • Barbon et al. (1990) Barbon R., et al., 1990, A&A, 237, 79
  • Benetti et al. (1995) Benetti S., et al., 1995, International Astronomical Union Circular, 6135, 1
  • Benetti et al. (2004) Benetti S., et al., 2004, MNRAS, 348, 261
  • Benetti et al. (2017) Benetti S., et al., 2017, The Astronomer’s Telegram, 11036, 1
  • Benn et al. (2008) Benn C., et al., 2008, in Ground-based and Airborne Instrumentation for Astronomy II. Edited by McLean, Ian S.; Casali, Mark M. Proceedings of the SPIE, Volume 7014, article id. 70146X, 12 pp. (2008).. , doi:10.1117/12.788694
  • Bertin & Arnouts (1996) Bertin E., et al., 1996, Astronomy and Astrophysics Supplement Series, 117, 393
  • Beutler & Li (2003) Beutler B., et al., 2003, International Astronomical Union Circular, 8197, 1
  • Bianco et al. (2011) Bianco F. B., et al., 2011, ApJ, 741, 20
  • Blanco et al. (1980) Blanco V. M., et al., 1980, International Astronomical Union Circular, 3556, 2
  • Blanco et al. (1986) Blanco V. M., et al., 1986, International Astronomical Union Circular, 4224, 1
  • Blondin & Berlind (2008) Blondin S., et al., 2008, Central Bureau Electronic Telegrams, 1198, 1
  • Blondin et al. (2012) Blondin S., et al., 2012, AJ, 143
  • Boehner et al. (2017) Boehner P., et al., 2017, MNRAS, 465, 2060
  • Botyánszki & Kasen (2017) Botyánszki J., et al., 2017, ApJ, 845
  • Botyánszki et al. (2018) Botyánszki J., et al., 2018, ApJ, 852
  • Branch et al. (1983) Branch D., et al., 1983, ApJ, 270, 123
  • Branch et al. (1993) Branch D., et al., 1993, AJ, 106, 2383
  • Breare et al. (1987) Breare J. M., et al., 1987, MNRAS, 227, 909
  • Brimacombe et al. (2011) Brimacombe J., et al., 2011, Central Bureau Electronic Telegrams, 2928, 1
  • Brimacombe et al. (2014a) Brimacombe J., et al., 2014a, The Astronomer’s Telegram, 6737, 1
  • Brimacombe et al. (2014b) Brimacombe J., et al., 2014b, The Astronomer’s Telegram, 6803, 1
  • Brimacombe et al. (2015) Brimacombe J., et al., 2015, The Astronomer’s Telegram, 6950, 1
  • Brimacombe et al. (2016) Brimacombe J., et al., 2016, The Astronomer’s Telegram, 8979, 1
  • Brimacombe et al. (2017) Brimacombe J., et al., 2017, The Astronomer’s Telegram, 10108, 1
  • Brimacombe et al. (2018) Brimacombe J., et al., 2018, The Astronomer’s Telegram, 11521, 1
  • Brown et al. (2012) Brown P. J., et al., 2012, ApJ, 749
  • Brown et al. (2014) Brown P. J., et al., 2014, Ap&SS, 354, 89
  • Brown et al. (2015) Brown P. J., et al., 2015, ApJ, 805, 74
  • Brown et al. (2016) Brown J. S., et al., 2016, The Astronomer’s Telegram, 9666, 1
  • Brown et al. (2018) Brown J. S., et al., 2018, The Astronomer’s Telegram, 11253, 1
  • Buckley et al. (2006) Buckley D. A. H., et al., 2006, in Society of Photo-Optical Instrumentation Engineers (SPIE) Conference Series. p. 62690A, doi:10.1117/12.673838
  • Bues et al. (1986) Bues I., et al., 1986, International Astronomical Union Circular, 4215, 2
  • Bureau et al. (1996) Bureau M., et al., 1996, ApJ, 463, 60
  • Burns et al. (2011) Burns C. R., et al., 2011, AJ, 141
  • Burns et al. (2014) Burns C. R., et al., 2014, ApJ, 789, 32
  • Busko et al. (1981) Busko I., et al., 1981, International Astronomical Union Circular, 3589, 2
  • Buta & Turner (1983) Buta R. J., et al., 1983, Publications of the Astronomical Society of the Pacific, 95, 72
  • Buzzoni et al. (1984) Buzzoni B., et al., 1984, The Messenger, 38, 9
  • Cacella et al. (2002) Cacella P., et al., 2002, International Astronomical Union Circular, 7847, 1
  • Canal et al. (2001) Canal R., et al., 2001, ApJ, 550, L53
  • Candia et al. (2003) Candia P., et al., 2003, Publications of the Astronomical Society of the Pacific, 115, 277
  • Cappellari et al. (2011) Cappellari M., et al., 2011, MNRAS, 413, 813
  • Cappellaro et al. (2001) Cappellaro E., et al., 2001, ApJ, 549, L215
  • Carnall (2017) Carnall A. C., 2017, preprint, p. arXiv:1705.05165 (arXiv:1705.05165)
  • Casper et al. (2013) Casper C., et al., 2013, Central Bureau Electronic Telegrams, 3588, 1
  • Catinella et al. (2005) Catinella B., et al., 2005, AJ, 130, 1037
  • Cenko et al. (2011) Cenko S. B., et al., 2011, The Astronomer’s Telegram, 3583, 1
  • Cenko et al. (2012) Cenko S. B., et al., 2012, Central Bureau Electronic Telegrams, 3014, 1
  • Cepa (2010) Cepa J., 2010, Astrophysics and Space Science Proceedings, 14, 15
  • Challis & Berlind (2009) Challis P., et al., 2009, Central Bureau Electronic Telegrams, 2025, 1
  • Chambers et al. (2016) Chambers K. C., et al., 2016, preprint, p. arXiv:1612.05560 (arXiv:1612.05560)
  • Childress et al. (2015) Childress M. J., et al., 2015, MNRAS, 454, 3816
  • Childress et al. (2016) Childress M. J., et al., 2016, Publications of the Astronomical Society of Australia, 33, e055
  • Chomiuk et al. (2016) Chomiuk L., et al., 2016, ApJ, 821, 119
  • Chornock et al. (2000) Chornock R., et al., 2000, International Astronomical Union Circular, 7463, 1
  • Christensen et al. (2003) Christensen L., et al., 2003, A&A, 401, 479
  • Colless et al. (2003) Colless M., et al., 2003, arXiv e-prints, pp astro–ph/0306581
  • Collobert et al. (2006) Collobert M., et al., 2006, MNRAS, 370, 1213
  • Contreras et al. (2010) Contreras C., et al., 2010, AJ, 139, 519
  • Corsini et al. (2003) Corsini E. M., et al., 2003, ApJ, 599, L29
  • Cortini et al. (2014) Cortini G., et al., 2014, Central Bureau Electronic Telegrams, 3911, 1
  • Cousins (1972) Cousins A. W. J., 1972, Information Bulletin on Variable Stars, 700, 1
  • Cox et al. (2010) Cox L., et al., 2010, Central Bureau Electronic Telegrams, 2612, 1
  • Cox et al. (2011) Cox L., et al., 2011, Central Bureau Electronic Telegrams, 2676, 1
  • Cragg et al. (1981) Cragg T., et al., 1981, International Astronomical Union Circular, 3583, 1
  • Cristiani et al. (1992) Cristiani S., et al., 1992, A&A, 259, 63
  • Cumming et al. (1994) Cumming R. J., et al., 1994, International Astronomical Union Circular, 5951, 1
  • D’Odorico (1990) D’Odorico S., 1990, The Messenger, 61, 51
  • D’Onofrio et al. (1995) D’Onofrio M., et al., 1995, A&A, 296, 319
  • Delgado et al. (2016) Delgado A., et al., 2016, Transient Name Server Discovery Report, 2016-485, 1
  • Di Stefano et al. (2011) Di Stefano R., et al., 2011, ApJ, 738, L1
  • Dilday et al. (2012) Dilday B., et al., 2012, Science, 337, 942
  • Dimitriadis et al. (2014) Dimitriadis G., et al., 2014, The Astronomer’s Telegram, 6749, 1
  • Dimitriadis et al. (2019) Dimitriadis G., et al., 2019, ApJ, 870, L14
  • Dong et al. (2015a) Dong S., et al., 2015a, MNRAS, 454, L61
  • Dong et al. (2015b) Dong S., et al., 2015b, The Astronomer’s Telegram, 7447, 1
  • Dong et al. (2018a) Dong S., et al., 2018a, MNRAS, 479, L70
  • Dong et al. (2018b) Dong S., et al., 2018b, The Astronomer’s Telegram, 11346, 1
  • Dopita et al. (2007) Dopita M., et al., 2007, Ap&SS, 310, 255
  • Dopita et al. (2010) Dopita M., et al., 2010, Ap&SS, 327, 245
  • Downes et al. (1993) Downes D., et al., 1993, ApJ, 414, L13
  • Drake et al. (2011a) Drake A. J., et al., 2011a, Central Bureau Electronic Telegrams, 2636, 1
  • Drake et al. (2011b) Drake A. J., et al., 2011b, Central Bureau Electronic Telegrams, 2703, 1
  • Drescher et al. (2012) Drescher C., et al., 2012, Central Bureau Electronic Telegrams, 3346, 1
  • Dressler et al. (2003) Dressler A., et al., 2003, International Astronomical Union Circular, 8198, 2
  • Dressler et al. (2011) Dressler A., et al., 2011, Publications of the Astronomical Society of the Pacific, 123, 288
  • Elias & Frogel (1983) Elias J. H., et al., 1983, ApJ, 268, 718
  • Elias-Rosa et al. (2005) Elias-Rosa N., et al., 2005, International Astronomical Union Circular, 8479, 3
  • Elias-Rosa et al. (2006) Elias-Rosa N., et al., 2006, MNRAS, 369, 1880
  • Epinat et al. (2008) Epinat B., et al., 2008, MNRAS, 388, 500
  • Evans et al. (1986) Evans R., et al., 1986, International Astronomical Union Circular, 4208, 1
  • Evans et al. (2003) Evans R., et al., 2003, International Astronomical Union Circular, 8171, 1
  • Evans et al. (2018) Evans D. W., et al., 2018, preprint, p. arXiv:1804.09368 (arXiv:1804.09368)
  • Everson et al. (2017) Everson R. W., et al., 2017, The Astronomer’s Telegram, 10518, 1
  • Faber et al. (2003) Faber S. M., et al., 2003, in Instrument Design and Performance for Optical/Infrared Ground-based Telescopes. Edited by Iye, Masanori; Moorwood, Alan F. M. Proceedings of the SPIE, Volume 4841, pp. 1657-1669 (2003).. pp 1657–1669, doi:10.1117/12.460346
  • Fabricant et al. (1998) Fabricant D., et al., 1998, Publications of the Astronomical Society of the Pacific, 110, 79
  • Falco et al. (1999) Falco E. E., et al., 1999, Publications of the Astronomical Society of the Pacific, 111, 438
  • Feast et al. (1986) Feast M. W., et al., 1986, International Astronomical Union Circular, 4210, 1
  • Filippenko et al. (1999) Filippenko A. V., et al., 1999, International Astronomical Union Circular, 7108, 2
  • Filippenko et al. (2007) Filippenko A. V., et al., 2007, Central Bureau Electronic Telegrams, 1101, 1
  • Fink et al. (2010) Fink M., et al., 2010, A&A, 514, A53
  • Firth et al. (2015) Firth R. E., et al., 2015, MNRAS, 446, 3895
  • Fitzpatrick (1999) Fitzpatrick E. L., 1999, Publications of the Astronomical Society of the Pacific, 111, 63
  • Flewelling et al. (2016) Flewelling H. A., et al., 2016, preprint, p. arXiv:1612.05243 (arXiv:1612.05243)
  • Folatelli et al. (2010a) Folatelli G., et al., 2010a, AJ, 139, 120
  • Folatelli et al. (2010b) Folatelli G., et al., 2010b, Central Bureau Electronic Telegrams, 2390, 1
  • Folatelli et al. (2013) Folatelli G., et al., 2013, ApJ, 773
  • Foley et al. (2013) Foley R. J., et al., 2013, ApJ, 767, 57
  • Foley et al. (2014) Foley R. J., et al., 2014, MNRAS, 443, 2887
  • Foley et al. (2015) Foley R. J., et al., 2015, ApJ, 798, L37
  • Foley et al. (2016) Foley R. J., et al., 2016, MNRAS, 461, 433
  • Foley et al. (2018) Foley R. J., et al., 2018, MNRAS, 475, 193
  • Ford et al. (1993) Ford C. H., et al., 1993, AJ, 106, 1101
  • Fossey et al. (2014) Fossey S. J., et al., 2014, Central Bureau Electronic Telegrams, 3792, 1
  • Freudling et al. (2013) Freudling W., et al., 2013, A&A, 559, A96
  • Friedman et al. (2015) Friedman A. S., et al., 2015, The Astrophysical Journal Supplement Series, 220
  • Frieman et al. (2006) Frieman J., et al., 2006, International Astronomical Union Circular, 8754, 1
  • Frohmaier et al. (2015) Frohmaier C., et al., 2015, The Astronomer’s Telegram, 7452, 1
  • Frye et al. (1972) Frye R., et al., 1972, International Astronomical Union Circular, 2398, 1
  • Fukugita et al. (1996) Fukugita M., et al., 1996, AJ, 111, 1748
  • Gagliano et al. (2016) Gagliano R., et al., 2016, Transient Name Server Discovery Report, 2016-761, 1
  • Gagliano et al. (2017) Gagliano R., et al., 2017, Transient Name Server Discovery Report, 2017-961, 1
  • Gaia Collaboration et al. (2016) Gaia Collaboration et al., 2016, A&A, 595, A1
  • Gaia Collaboration et al. (2018) Gaia Collaboration et al., 2018, preprint, p. arXiv:1804.09365 (arXiv:1804.09365)
  • Galbany et al. (2014) Galbany L., et al., 2014, A&A, 572, A38
  • Galbany et al. (2016) Galbany L., et al., 2016, MNRAS, 457, 525
  • Gall et al. (2018) Gall C., et al., 2018, A&A, 611, A58
  • Ganeshalingam et al. (2010) Ganeshalingam M., et al., 2010, The Astrophysical Journal Supplement Series, 190, 418
  • Gao et al. (2015) Gao J., et al., 2015, ApJ, 807, L26
  • Garnavich et al. (2004) Garnavich P. M., et al., 2004, ApJ, 613, 1120
  • Garradd et al. (1996) Garradd G. J., et al., 1996, International Astronomical Union Circular, 6380, 1
  • Gaskell et al. (1989) Gaskell C. M., et al., 1989, International Astronomical Union Circular, 4761, 2
  • Gerardy & Fesen (1999) Gerardy C., et al., 1999, International Astronomical Union Circular, 7158, 2
  • Gilfanov & Bogdán (2010) Gilfanov M., et al., 2010, Nature, 463, 924
  • Gilmore (1991) Gilmore A. C., 1991, International Astronomical Union Circular, 5309, 3
  • Ginsburg et al. (2019) Ginsburg A., et al., 2019, arXiv e-prints, p. arXiv:1901.04520
  • Giovanelli et al. (1997) Giovanelli R., et al., 1997, AJ, 113, 22
  • Gomez et al. (1996) Gomez G., et al., 1996, AJ, 112, 2094
  • Gonzalez et al. (2004) Gonzalez S., et al., 2004, International Astronomical Union Circular, 8409, 2
  • Goobar et al. (2014) Goobar A., et al., 2014, ApJ, 784
  • Graham et al. (1978) Graham D. A., et al., 1978, A&A, 70, L69
  • Graham et al. (1998) Graham A. W., et al., 1998, Astronomy and Astrophysics Supplement Series, 133, 325
  • Graham et al. (2017) Graham M. L., et al., 2017, MNRAS, 472, 3437
  • Graham et al. (2019) Graham M. L., et al., 2019, ApJ, 871, 62
  • Grogin et al. (1998) Grogin N. A., et al., 1998, The Astrophysical Journal Supplement Series, 119, 277
  • Guillochon et al. (2017) Guillochon J., et al., 2017, ApJ, 835
  • Guthrie & Napier (1996) Guthrie B. N. G., et al., 1996, A&A, 310, 353
  • Gutiérrez et al. (2016) Gutiérrez C. P., et al., 2016, A&A, 590, A5
  • Halevi et al. (2016) Halevi G., et al., 2016, The Astronomer’s Telegram, 9309, 1
  • Hamuy et al. (1991) Hamuy M., et al., 1991, AJ, 102, 208
  • Hamuy et al. (1996) Hamuy M., et al., 1996, AJ, 112, 2408
  • Hamuy et al. (2006) Hamuy M., et al., 2006, Publications of the Astronomical Society of the Pacific, 118, 2
  • Harutyunyan et al. (2009) Harutyunyan A., et al., 2009, Central Bureau Electronic Telegrams, 1768, 1
  • Hayden et al. (2010) Hayden B. T., et al., 2010, ApJ, 722, 1691
  • Heraudeau et al. (1994) Heraudeau P., et al., 1994, International Astronomical Union Circular, 5952, 3
  • Herbig et al. (1972) Herbig G. H., et al., 1972, International Astronomical Union Circular, 2407, 1
  • Hernandez et al. (2000) Hernandez M., et al., 2000, MNRAS, 319, 223
  • Hicken et al. (2007) Hicken M., et al., 2007, ApJ, 669, L17
  • Hicken et al. (2009) Hicken M., et al., 2009, ApJ, 700, 331
  • Hicken et al. (2012) Hicken M., et al., 2012, The Astrophysical Journal Supplement Series, 200
  • Ho et al. (2001) Ho W. C. G., et al., 2001, Publications of the Astronomical Society of the Pacific, 113, 1349
  • Höflich et al. (2004) Höflich P., et al., 2004, ApJ, 617, 1258
  • Holberg et al. (1991) Holberg J., et al., 1991, International Astronomical Union Circular, 5270, 3
  • Holmbo et al. (2018) Holmbo S., et al., 2018, arXiv e-prints, p. arXiv:1809.01359
  • Holoien et al. (2014) Holoien T. W. S., et al., 2014, The Astronomer’s Telegram, 6637, 1
  • Holoien et al. (2017a) Holoien T. W. S., et al., 2017a, MNRAS, 464, 2672
  • Holoien et al. (2017b) Holoien T. W. S., et al., 2017b, MNRAS, 467, 1098
  • Holoien et al. (2017c) Holoien T. W. S., et al., 2017c, MNRAS, 471, 4966
  • Holoien et al. (2019) Holoien T. W. S., et al., 2019, MNRAS, p. 93
  • Holtzman et al. (2008) Holtzman J. A., et al., 2008, AJ, 136, 2306
  • Hook et al. (2004) Hook I. M., et al., 2004, Publications of the Astronomical Society of the Pacific, 116, 425
  • Hosseinzadeh et al. (2017a) Hosseinzadeh G., et al., 2017a, The Astronomer’s Telegram, 10164, 1
  • Hosseinzadeh et al. (2017b) Hosseinzadeh G., et al., 2017b, The Astronomer’s Telegram, 10639, 1
  • Howell et al. (2006) Howell D. A., et al., 2006, Nature, 443, 308
  • Howerton et al. (2011) Howerton S., et al., 2011, Central Bureau Electronic Telegrams, 2658, 1
  • Howerton et al. (2013) Howerton S., et al., 2013, Central Bureau Electronic Telegrams, 3533, 1
  • Hoyle & Fowler (1960) Hoyle F., et al., 1960, ApJ, 132, 565
  • Huchra et al. (1999) Huchra J. P., et al., 1999, The Astrophysical Journal Supplement Series, 121, 287
  • Hurst et al. (1998) Hurst G. M., et al., 1998, International Astronomical Union Circular, 6875, 1
  • Hutchings & Li (2002) Hutchings D., et al., 2002, International Astronomical Union Circular, 7918, 1
  • Iben & Tutukov (1984) Iben I., et al., 1984, The Astrophysical Journal Supplement Series, 54, 335
  • Iijima et al. (1994) Iijima T., et al., 1994, International Astronomical Union Circular, 6108, 1
  • Inserra et al. (2013) Inserra C., et al., 2013, ApJ, 770
  • Itagaki et al. (2014) Itagaki K., et al., 2014, Central Bureau Electronic Telegrams, 3792, 2
  • Jha et al. (1999) Jha S., et al., 1999, The Astrophysical Journal Supplement Series, 125, 73
  • Jha et al. (2006) Jha S., et al., 2006, AJ, 131, 527
  • Jha et al. (2017) Jha S. W., et al., 2017, The Astronomer’s Telegram, 10261, 1
  • Jones et al. (2006) Jones D. H., et al., 2006, MNRAS, 369, 25
  • Jones et al. (2009) Jones D. H., et al., 2009, MNRAS, 399, 683
  • Jones et al. (2017) Jones S., et al., 2017, The Astronomer’s Telegram, 11092, 1
  • Jorden (1990) Jorden P. R., 1990, in Proc. SPIE Vol. 1235, p. 790-798, Instrumentation in Astronomy VII, David L. Crawford; Ed.. pp 790–798, doi:10.1117/12.19163
  • Kandrashoff et al. (2012) Kandrashoff M., et al., 2012, Central Bureau Electronic Telegrams, 3111, 1
  • Karamehmetoglu et al. (2015) Karamehmetoglu E., et al., 2015, The Astronomer’s Telegram, 7476, 1
  • Kasen (2010) Kasen D., 2010, ApJ, 708, 1025
  • Kashikawa et al. (2002) Kashikawa N., et al., 2002, Publications of the Astronomical Society of Japan, 54, 819
  • Katz & Dong (2012) Katz B., et al., 2012, preprint, p. arXiv:1211.4584 (arXiv:1211.4584)
  • Kawakita et al. (2002) Kawakita H., et al., 2002, International Astronomical Union Circular, 7848, 2
  • Kelson (2003) Kelson D. D., 2003, Publications of the Astronomical Society of the Pacific, 115, 688
  • Kelson et al. (2000) Kelson D. D., et al., 2000, ApJ, 531, 159
  • Kent et al. (2008) Kent B. R., et al., 2008, AJ, 136, 713
  • Khan et al. (2011) Khan R., et al., 2011, ApJ, 726, 106
  • Kim et al. (2013) Kim H., et al., 2013, Central Bureau Electronic Telegrams, 3743, 1
  • Kirshner & Oke (1975) Kirshner R. P., et al., 1975, ApJ, 200, 574
  • Kirshner et al. (1991) Kirshner R., et al., 1991, International Astronomical Union Circular, 5403, 2
  • Kirshner et al. (1993) Kirshner R. P., et al., 1993, ApJ, 415, 589
  • Kiyota et al. (2014a) Kiyota S., et al., 2014a, The Astronomer’s Telegram, 6594, 1
  • Kiyota et al. (2014b) Kiyota S., et al., 2014b, The Astronomer’s Telegram, 6683, 1
  • Kiyota et al. (2014c) Kiyota S., et al., 2014c, The Astronomer’s Telegram, 6802, 1
  • Kiyota et al. (2014d) Kiyota S., et al., 2014d, The Astronomer’s Telegram, 6809, 1
  • Kleiser et al. (2009) Kleiser I., et al., 2009, Central Bureau Electronic Telegrams, 1918, 1
  • Klotz et al. (2012) Klotz A., et al., 2012, Central Bureau Electronic Telegrams, 3277, 2
  • Kollmeier et al. (2019) Kollmeier J. A., et al., 2019, arXiv e-prints, p. arXiv:1902.02251
  • Koribalski et al. (2004) Koribalski B. S., et al., 2004, AJ, 128, 16
  • Kosai et al. (1991) Kosai H., et al., 1991, International Astronomical Union Circular, 5400, 1
  • Kotak et al. (2003a) Kotak R., et al., 2003a, International Astronomical Union Circular, 8099, 1
  • Kotak et al. (2003b) Kotak R., et al., 2003b, International Astronomical Union Circular, 8122, 3
  • Kotak et al. (2005) Kotak R., et al., 2005, A&A, 436, 1021
  • Kowal (1972) Kowal C. T., 1972, International Astronomical Union Circular, 2405, 1
  • Kowalski et al. (2008) Kowalski M., et al., 2008, ApJ, 686
  • Krisciunas et al. (2000) Krisciunas K., et al., 2000, ApJ, 539, 658
  • Krisciunas et al. (2003) Krisciunas K., et al., 2003, AJ, 125, 166
  • Krisciunas et al. (2004) Krisciunas K., et al., 2004, AJ, 128, 3034
  • Krisciunas et al. (2007) Krisciunas K., et al., 2007, AJ, 133, 58
  • Krisciunas et al. (2009) Krisciunas K., et al., 2009, AJ, 138, 1584
  • Krisciunas et al. (2017) Krisciunas K., et al., 2017, AJ, 154, 211
  • Kromer et al. (2010) Kromer M., et al., 2010, ApJ, 719, 1067
  • Krumm & Salpeter (1980) Krumm N., et al., 1980, AJ, 85, 1312
  • Lair et al. (2006) Lair J. C., et al., 2006, AJ, 132, 2024
  • Lanz et al. (2005) Lanz T., et al., 2005, ApJ, 619, 517
  • Lauberts & Valentijn (1989) Lauberts A., et al., 1989, The surface photometry catalogue of the ESO-Uppsala galaxies
  • Lavery (1989) Lavery R. J., 1989, International Astronomical Union Circular, 4743, 2
  • Leadbeater (2016) Leadbeater R., 2016, Transient Name Server Classification Report, 2016-793, 1
  • Leadbeater (2018) Leadbeater R., 2018, Transient Name Server Classification Report, 2018-159, 1
  • Lee et al. (1972) Lee T. A., et al., 1972, ApJ, 177, L59
  • Leibendgut et al. (1993) Leibendgut B., et al., 1993, AJ, 105, 301
  • Leloudas et al. (2009) Leloudas G., et al., 2009, A&A, 505, 265
  • Leonard (2007) Leonard D. C., 2007, ApJ, 670, 1275
  • Leonard & Filippenko (2001) Leonard D. C., et al., 2001, Publications of the Astronomical Society of the Pacific, 113, 920
  • Li et al. (1999) Li W. D., et al., 1999, AJ, 117, 2709
  • Li et al. (2001) Li W., et al., 2001, Publications of the Astronomical Society of the Pacific, 113, 1178
  • Li et al. (2003a) Li W., et al., 2003a, Publications of the Astronomical Society of the Pacific, 115, 453
  • Li et al. (2003b) Li W., et al., 2003b, International Astronomical Union Circular, 8245, 1
  • Liller & Buta (1992) Liller W., et al., 1992, International Astronomical Union Circular, 5431, 3
  • Liller et al. (1992) Liller W., et al., 1992, International Astronomical Union Circular, 5428, 1
  • Lira et al. (1998) Lira P., et al., 1998, AJ, 115, 234
  • Liu et al. (2012) Liu Z. W., et al., 2012, A&A, 548, A2
  • Liu et al. (2013a) Liu Z.-W., et al., 2013a, ApJ, 774, 37
  • Liu et al. (2013b) Liu Z.-W., et al., 2013b, ApJ, 778, 121
  • Livio & Mazzali (2018) Livio M., et al., 2018, Phys. Rep., 736, 1
  • Livne (1990) Livne E., 1990, ApJ, 354, L53
  • Livne & Arnett (1995) Livne E., et al., 1995, ApJ, 452, 62
  • Longhetti et al. (1998) Longhetti M., et al., 1998, Astronomy and Astrophysics Supplement Series, 130, 251
  • Lundqvist et al. (2013) Lundqvist P., et al., 2013, MNRAS, 435, 329
  • Lundqvist et al. (2015) Lundqvist P., et al., 2015, A&A, 577, A39
  • Maeda et al. (2009) Maeda K., et al., 2009, ApJ, 690, 1745
  • Maguire et al. (2014) Maguire K., et al., 2014, MNRAS, 444, 3258
  • Maguire et al. (2016) Maguire K., et al., 2016, MNRAS, 457, 3254
  • Maguire et al. (2018) Maguire K., et al., 2018, MNRAS, 477, 3567
  • Malesani et al. (2018) Malesani D., et al., 2018, The Astronomer’s Telegram, 11516, 1
  • Maoz et al. (2014) Maoz D., et al., 2014, Annual Review of Astronomy and Astrophysics, 52, 107
  • Marietta et al. (2000) Marietta E., et al., 2000, The Astrophysical Journal Supplement Series, 128, 615
  • Marion et al. (2012) Marion G. H., et al., 2012, Central Bureau Electronic Telegrams, 3146, 2
  • Marshall et al. (2008) Marshall J. L., et al., 2008, in Ground-based and Airborne Instrumentation for Astronomy II. Edited by McLean, Ian S.; Casali, Mark M. Proceedings of the SPIE, Volume 7014, article id. 701454, 10 pp. (2008).. , doi:10.1117/12.789972
  • Martin & Biggs (2004) Martin R., et al., 2004, International Astronomical Union Circular, 8282, 1
  • Martin et al. (2005) Martin R., et al., 2005, International Astronomical Union Circular, 8490, 1
  • Matheson et al. (2002) Matheson T., et al., 2002, International Astronomical Union Circular, 7903, 2
  • Matheson et al. (2008) Matheson T., et al., 2008, AJ, 135, 1598
  • Mathewson et al. (1992) Mathewson D. S., et al., 1992, The Astrophysical Journal Supplement Series, 81, 413
  • Mattila et al. (2005) Mattila S., et al., 2005, A&A, 443, 649
  • Mattila et al. (2016) Mattila S., et al., 2016, The Astronomer’s Telegram, 8992, 1
  • Maury et al. (1990) Maury A., et al., 1990, International Astronomical Union Circular, 5039, 1
  • Maza et al. (2010) Maza J., et al., 2010, Central Bureau Electronic Telegrams, 2388, 1
  • Mazzali et al. (2015) Mazzali P. A., et al., 2015, MNRAS, 450, 2631
  • McCully et al. (2014) McCully C., et al., 2014, Nature, 512, 54
  • Menzies et al. (1991) Menzies J., et al., 1991, International Astronomical Union Circular, 5246, 2
  • Meyer et al. (2004) Meyer M. J., et al., 2004, MNRAS, 350, 1195
  • Mikuz (1994) Mikuz H., 1994, International Astronomical Union Circular, 5958, 3
  • Mikuz (1995) Mikuz H., 1995, International Astronomical Union Circular, 6166, 2
  • Misra et al. (2005) Misra K., et al., 2005, MNRAS, 360, 662
  • Modjaz et al. (2005a) Modjaz M., et al., 2005a, Central Bureau Electronic Telegrams, 160, 1
  • Modjaz et al. (2005b) Modjaz M., et al., 2005b, International Astronomical Union Circular, 8491, 2
  • Molinari et al. (1999) Molinari E., et al., 1999, in Looking Deep in the Southern Sky, Proceedings of the ESO/Australia Workshop held at Sydney, Australia, 10-12 December 1997. Edited by Faffaella Morganti and Warrick J. Couch. Berlin: Springer- Verlag, 1999. p. 157.. p. 157
  • Monard & Africa (2010) Monard L. A. G., et al., 2010, Central Bureau Electronic Telegrams, 2434, 1
  • Monard et al. (2001) Monard A. G., et al., 2001, International Astronomical Union Circular, 7720, 1
  • Monard et al. (2007) Monard L. A. G., et al., 2007, Central Bureau Electronic Telegrams, 1100, 1
  • Monard et al. (2011) Monard L. A. G., et al., 2011, Central Bureau Electronic Telegrams, 2635, 1
  • Monard et al. (2015) Monard L. A. G., et al., 2015, Central Bureau Electronic Telegrams, 4081, 1
  • Moorwood et al. (1998) Moorwood A., et al., 1998, The Messenger, 91, 9
  • Morrell & Shappee (2016) Morrell N., et al., 2016, The Astronomer’s Telegram, 9170, 1
  • Morrell et al. (2007) Morrell N., et al., 2007, Central Bureau Electronic Telegrams, 1131, 1
  • Morrell et al. (2014a) Morrell N., et al., 2014a, The Astronomer’s Telegram, 6508, 1
  • Morrell et al. (2014b) Morrell N., et al., 2014b, The Astronomer’s Telegram, 6765, 1
  • Morrell et al. (2015) Morrell N., et al., 2015, The Astronomer’s Telegram, 6988, 1
  • Motohara et al. (2002) Motohara K., et al., 2002, Publications of the Astronomical Society of Japan, 54, 315
  • Motohara et al. (2006) Motohara K., et al., 2006, ApJ, 652, L101
  • Munari et al. (2013) Munari U., et al., 2013, New Astron., 20, 30
  • Munch et al. (1986) Munch G., et al., 1986, International Astronomical Union Circular, 4183, 1
  • Nakano & Itagaki (2007) Nakano S., et al., 2007, Central Bureau Electronic Telegrams, 863, 1
  • Nakano et al. (1995) Nakano S., et al., 1995, International Astronomical Union Circular, 6134, 1
  • Nakano et al. (2003) Nakano S., et al., 2003, International Astronomical Union Circular, 8097, 1
  • Nakano et al. (2005) Nakano S., et al., 2005, International Astronomical Union Circular, 8475, 1
  • Nakano et al. (2008) Nakano S., et al., 2008, Central Bureau Electronic Telegrams, 1193, 1
  • Nakano et al. (2011) Nakano S., et al., 2011, Central Bureau Electronic Telegrams, 2783, 1
  • Nakano et al. (2012) Nakano S., et al., 2012, Central Bureau Electronic Telegrams, 3209, 1
  • Nakano et al. (2015) Nakano S., et al., 2015, Central Bureau Electronic Telegrams, 4106, 1
  • Nicolas et al. (2014) Nicolas J., et al., 2014, The Astronomer’s Telegram, 6500, 1
  • Nishiyama et al. (2012) Nishiyama K., et al., 2012, Central Bureau Electronic Telegrams, 3349, 1
  • Noguchi et al. (2011a) Noguchi T., et al., 2011a, Central Bureau Electronic Telegrams, 2940, 1
  • Noguchi et al. (2011b) Noguchi T., et al., 2011b, Central Bureau Electronic Telegrams, 2943, 1
  • Nomoto (1982) Nomoto K., 1982, ApJ, 253, 798
  • Nucita et al. (2017) Nucita A. A., et al., 2017, ApJ, 850, 111
  • Nugent et al. (2011a) Nugent P. E., et al., 2011a, Nature, 480, 344
  • Nugent et al. (2011b) Nugent P., et al., 2011b, The Astronomer’s Telegram, 3581, 1
  • Nyholm et al. (2017) Nyholm A., et al., 2017, The Astronomer’s Telegram, 10131, 1
  • Ochner et al. (2016) Ochner P., et al., 2016, The Astronomer’s Telegram, 9018, 1
  • Ogando et al. (2008) Ogando R. L. C., et al., 2008, AJ, 135, 2424
  • Oke & Gunn (1982) Oke J. B., et al., 1982, Publications of the Astronomical Society of the Pacific, 94, 586
  • Oke et al. (1995) Oke J. B., et al., 1995, Publications of the Astronomical Society of the Pacific, 107, 375
  • Olszewski (1982) Olszewski E. W., 1982, Information Bulletin on Variable Stars, 2065, 1
  • Osmer et al. (1972) Osmer P. S., et al., 1972, Nature Physical Science, 238, 21
  • Pakmor et al. (2012) Pakmor R., et al., 2012, ApJ, 747, L10
  • Pan et al. (2012) Pan K.-C., et al., 2012, ApJ, 750, 151
  • Pan et al. (2015) Pan Y. C., et al., 2015, MNRAS, 452, 4307
  • Pan et al. (2017) Pan Y. C., et al., 2017, The Astronomer’s Telegram, 10225, 1
  • Parker (2016) Parker S., 2016, Transient Name Server Discovery Report, 2016-304, 1
  • Parker (2017) Parker S., 2017, Transient Name Server Discovery Report, 2017-700, 1
  • Parker et al. (2013a) Parker S., et al., 2013a, Central Bureau Electronic Telegrams, 3416, 1
  • Parker et al. (2013b) Parker P., et al., 2013b, Central Bureau Electronic Telegrams, 3539, 1
  • Pastorello et al. (2007) Pastorello A., et al., 2007, MNRAS, 376, 1301
  • Paturel et al. (2003) Paturel G., et al., 2003, A&A, 412, 57
  • Pejcha et al. (2013) Pejcha O., et al., 2013, MNRAS, 435, 943
  • Pessa et al. (2018) Pessa I., et al., 2018, preprint, p. arXiv:1803.05005 (arXiv:1803.05005)
  • Petrushevska et al. (2016) Petrushevska T., et al., 2016, The Astronomer’s Telegram, 9049, 1
  • Phillips et al. (1987) Phillips M. M., et al., 1987, Publications of the Astronomical Society of the Pacific, 99, 592
  • Phillips et al. (2006) Phillips M. M., et al., 2006, AJ, 131, 2615
  • Phillips et al. (2007) Phillips M. M., et al., 2007, Publications of the Astronomical Society of the Pacific, 119, 360
  • Phillips et al. (2013) Phillips M. M., et al., 2013, ApJ, 779, 38
  • Phillips et al. (2019) Phillips M. M., et al., 2019, Publications of the Astronomical Society of the Pacific, 131, 014001
  • Pignata et al. (2004) Pignata G., et al., 2004, MNRAS, 355, 178
  • Pignata et al. (2008) Pignata G., et al., 2008, MNRAS, 388, 971
  • Pignata et al. (2009) Pignata G., et al., 2009, Central Bureau Electronic Telegrams, 2022, 1
  • Pignata et al. (2010) Pignata G., et al., 2010, Central Bureau Electronic Telegrams, 2344, 1
  • Pisano et al. (2011) Pisano D. J., et al., 2011, The Astrophysical Journal Supplement Series, 197, 28
  • Pogge et al. (2010) Pogge R. W., et al., 2010, in Proceedings of the SPIE, Volume 7735, id. 77350A (2010).. , doi:10.1117/12.857215
  • Pollas & Klotz (2007) Pollas C., et al., 2007, Central Bureau Electronic Telegrams, 1121, 1
  • Pollas et al. (1989) Pollas C., et al., 1989, International Astronomical Union Circular, 4742, 2
  • Ponticello et al. (2006) Ponticello N. J., et al., 2006, International Astronomical Union Circular, 8667, 1
  • Prentice & Ashall (2017) Prentice S., et al., 2017, Transient Name Server Classification Report, 2017-978, 1
  • Press et al. (1992) Press W. H., et al., 1992, Numerical recipes in FORTRAN. The art of scientific computing
  • Prieto & Morrell (2011) Prieto J. L., et al., 2011, Central Bureau Electronic Telegrams, 2613, 1
  • Prieto et al. (2006) Prieto J. L., et al., 2006, Central Bureau Electronic Telegrams, 651, 1
  • Prieto et al. (2010) Prieto J. L., et al., 2010, Central Bureau Electronic Telegrams, 2453, 1
  • Przybylski (1972) Przybylski A., 1972, International Astronomical Union Circular, 2434, 1
  • Puckett et al. (2009) Puckett T., et al., 2009, Central Bureau Electronic Telegrams, 1762, 1
  • Pugh & Li (2005) Pugh H., et al., 2005, Central Bureau Electronic Telegrams, 158, 1
  • Quimby et al. (2005) Quimby R., et al., 2005, International Astronomical Union Circular, 8625, 2
  • Quimby et al. (2006a) Quimby R., et al., 2006a, Central Bureau Electronic Telegrams, 393, 1
  • Quimby et al. (2006b) Quimby R., et al., 2006b, International Astronomical Union Circular, 8723, 2
  • Rand (1995) Rand R. J., 1995, AJ, 109, 2444
  • Rebassa-Mansergas et al. (2018) Rebassa-Mansergas A., et al., 2018, preprint, p. arXiv:1809.07158 (arXiv:1809.07158)
  • Rhee & van Albada (1996) Rhee M. H., et al., 1996, Astronomy and Astrophysics Supplement Series, 115, 407
  • Richmond & Smith (2012) Richmond M. W., et al., 2012, Journal of the American Association of Variable Star Observers (JAAVSO), 40, 872
  • Richmond et al. (1995) Richmond M. W., et al., 1995, AJ, 109, 2121
  • Riello et al. (2002) Riello M., et al., 2002, International Astronomical Union Circular, 7919, 2
  • Riello et al. (2018) Riello M., et al., 2018, preprint, p. arXiv:1804.09367 (arXiv:1804.09367)
  • Riess et al. (1999) Riess A. G., et al., 1999, AJ, 117, 707
  • Riess et al. (2005) Riess A. G., et al., 2005, ApJ, 627, 579
  • Romero-Canizales et al. (2014) Romero-Canizales C., et al., 2014, The Astronomer’s Telegram, 6618, 1
  • Röpke et al. (2012) Röpke F. K., et al., 2012, ApJ, 750, L19
  • Rothberg & Joseph (2006) Rothberg B., et al., 2006, AJ, 131, 185
  • Rudy et al. (2015) Rudy R. J., et al., 2015, The Astronomer’s Telegram, 7825, 1
  • Sahu et al. (2008) Sahu D. K., et al., 2008, ApJ, 680, 580
  • Sako et al. (2014) Sako M., et al., 2014, preprint, (arXiv:1401.3317)
  • Salgado et al. (2007) Salgado F., et al., 2007, Central Bureau Electronic Telegrams, 865, 1
  • Salvo et al. (2001) Salvo M. E., et al., 2001, MNRAS, 321, 254
  • Salvo et al. (2006) Salvo M., et al., 2006, Central Bureau Electronic Telegrams, 557, 1
  • Sand et al. (2017) Sand D. J., et al., 2017, The Astronomer’s Telegram, 10569, 1
  • Sand et al. (2018a) Sand D. J., et al., 2018a, preprint, p. arXiv:1804.03666 (arXiv:1804.03666)
  • Sand et al. (2018b) Sand D., et al., 2018b, The Astronomer’s Telegram, 11328, 1
  • Sand et al. (2018c) Sand D., et al., 2018c, The Astronomer’s Telegram, 11330, 1
  • Sand et al. (2018d) Sand D., et al., 2018d, The Astronomer’s Telegram, 11371, 1
  • Scalzo et al. (2010) Scalzo R. A., et al., 2010, ApJ, 713, 1073
  • Scalzo et al. (2019) Scalzo R. A., et al., 2019, MNRAS, 483, 628
  • Schaefer (1987) Schaefer B., 1987, International Astronomical Union Circular, 4421, 1
  • Schlegel et al. (1998) Schlegel D. J., et al., 1998, ApJ, 500, 525
  • Schmidt et al. (1994) Schmidt B. P., et al., 1994, ApJ, 434, L19
  • Schneider et al. (1990) Schneider S. E., et al., 1990, The Astrophysical Journal Supplement Series, 72, 245
  • Schneider et al. (1992) Schneider S. E., et al., 1992, The Astrophysical Journal Supplement Series, 81, 5
  • Schwartz & Holvorcem (2003) Schwartz M., et al., 2003, International Astronomical Union Circular, 8121, 1
  • Serduke et al. (2005) Serduke F. J. D., et al., 2005, Central Bureau Electronic Telegrams, 269, 1
  • Shappee & Thompson (2013) Shappee B. J., et al., 2013, ApJ, 766
  • Shappee et al. (2013a) Shappee B. J., et al., 2013a, ApJ, 762
  • Shappee et al. (2013b) Shappee B. J., et al., 2013b, ApJ, 765, 150
  • Shappee et al. (2014a) Shappee B. J., et al., 2014a, ApJ, 788, 48
  • Shappee et al. (2014b) Shappee B. J., et al., 2014b, The Astronomer’s Telegram, 6812, 1
  • Shappee et al. (2015) Shappee B. J., et al., 2015, The Astronomer’s Telegram, 6882, 1
  • Shappee et al. (2017) Shappee B. J., et al., 2017, ApJ, 841, 48
  • Shappee et al. (2018) Shappee B. J., et al., 2018, ApJ, 855, 6
  • Sheinis et al. (2002) Sheinis A. I., et al., 2002, Publications of the Astronomical Society of the Pacific, 114, 851
  • Shen et al. (2018) Shen K. J., et al., 2018, preprint, p. arXiv:1804.11163 (arXiv:1804.11163)
  • Silverman et al. (2011) Silverman J. M., et al., 2011, MNRAS, 410, 585
  • Silverman et al. (2012) Silverman J. M., et al., 2012, MNRAS, 425, 1789
  • Silverman et al. (2013) Silverman J. M., et al., 2013, MNRAS, 430, 1030
  • Smartt et al. (2002) Smartt S. J., et al., 2002, International Astronomical Union Circular, 7961, 2
  • Smartt et al. (2015a) Smartt S. J., et al., 2015a, A&A, 579
  • Smartt et al. (2015b) Smartt S. J., et al., 2015b, A&A, 579, A40
  • Smith et al. (2000) Smith R. J., et al., 2000, MNRAS, 313, 469
  • Smoker et al. (2000) Smoker J. V., et al., 2000, A&A, 361, 19
  • Sollerman et al. (2001) Sollerman J., et al., 2001, International Astronomical Union Circular, 7723, 2
  • Sollerman et al. (2004) Sollerman J., et al., 2004, A&A, 428, 555
  • Springob et al. (2014) Springob C. M., et al., 2014, MNRAS, 445, 2677
  • Spyromilio et al. (2004) Spyromilio J., et al., 2004, A&A, 426, 547
  • Srivastav et al. (2016) Srivastav S., et al., 2016, MNRAS, 457, 1000
  • Stanek (2018) Stanek K. Z., 2018, Transient Name Server Discovery Report, 2018-234, 1
  • Stanek et al. (2014) Stanek K. Z., et al., 2014, The Astronomer’s Telegram, 6830, 1
  • Stanishev & Pursimo (2008) Stanishev V., et al., 2008, Central Bureau Electronic Telegrams, 1232, 1
  • Stanishev et al. (2007) Stanishev V., et al., 2007, A&A, 469, 645
  • Stone et al. (2018) Stone G., et al., 2018, The Astronomer’s Telegram, 11343, 1
  • Strauss et al. (1992) Strauss M. A., et al., 1992, The Astrophysical Journal Supplement Series, 83, 29
  • Stritzinger (2010) Stritzinger M., 2010, Central Bureau Electronic Telegrams, 2346, 1
  • Stritzinger et al. (2006) Stritzinger M., et al., 2006, A&A, 460, 793
  • Stritzinger et al. (2010) Stritzinger M., et al., 2010, AJ, 140, 2036
  • Stritzinger et al. (2011) Stritzinger M. D., et al., 2011, AJ, 142
  • Stritzinger et al. (2015) Stritzinger M. D., et al., 2015, A&A, 573, A2
  • Strohmayer et al. (1996) Strohmayer T., et al., 1996, International Astronomical Union Circular, 6484, 1
  • Suntzeff et al. (1989) Suntzeff N., et al., 1989, International Astronomical Union Circular, 4728, 1
  • Suntzeff et al. (1992) Suntzeff N., et al., 1992, International Astronomical Union Circular, 5432, 2
  • Suntzeff et al. (1996) Suntzeff N. B., et al., 1996, International Astronomical Union Circular, 6381, 1
  • Suntzeff et al. (1999) Suntzeff N. B., et al., 1999, AJ, 117, 1175
  • Suntzeff et al. (2004) Suntzeff N., et al., 2004, International Astronomical Union Circular, 8283, 1
  • Szabó et al. (2003) Szabó G. M., et al., 2003, A&A, 408, 915
  • Taam (1980) Taam R. E., 1980, ApJ, 237, 142
  • Tartaglia et al. (2017a) Tartaglia L., et al., 2017a, The Astronomer’s Telegram, 10158, 1
  • Tartaglia et al. (2017b) Tartaglia L., et al., 2017b, The Astronomer’s Telegram, 10260, 1
  • Tartaglia et al. (2017c) Tartaglia L., et al., 2017c, The Astronomer’s Telegram, 10439, 1
  • Tartaglia et al. (2017d) Tartaglia L., et al., 2017d, The Astronomer’s Telegram, 10629, 1
  • Taubenberger et al. (2013a) Taubenberger S., et al., 2013a, MNRAS, 432, 3117
  • Taubenberger et al. (2013b) Taubenberger S., et al., 2013b, ApJ, 775, L43
  • Taubenberger et al. (2015) Taubenberger S., et al., 2015, MNRAS, 448, L48
  • Terreran et al. (2016) Terreran G., et al., 2016, The Astronomer’s Telegram, 9403, 1
  • The Astropy Collaboration et al. (2018) The Astropy Collaboration et al., 2018, preprint, p. arXiv:1801.02634 (arXiv:1801.02634)
  • Theureau et al. (1998) Theureau G., et al., 1998, Astronomy and Astrophysics Supplement Series, 130, 333
  • Theureau et al. (2007) Theureau G., et al., 2007, A&A, 465, 71
  • Thompson (2011) Thompson T. A., 2011, ApJ, 741, 82
  • Tinella (2016) Tinella V., 2016, Transient Name Server Discovery Report, 2016-305, 1
  • Tonry et al. (2012) Tonry J. L., et al., 2012, ApJ, 750, 99
  • Tonry et al. (2016) Tonry J., et al., 2016, Transient Name Server Discovery Report, 2016-583, 1
  • Tonry et al. (2017a) Tonry J., et al., 2017a, Transient Name Server Discovery Report, 2017-1371, 1
  • Tonry et al. (2017b) Tonry J., et al., 2017b, Transient Name Server Discovery Report, 2017-1431, 1
  • Tonry et al. (2017c) Tonry J., et al., 2017c, Transient Name Server Discovery Report, 2017-361, 1
  • Tonry et al. (2017d) Tonry J., et al., 2017d, Transient Name Server Discovery Report, 2017-860, 1
  • Toth & Szabó (2000) Toth I., et al., 2000, A&A, 361, 63
  • Treffers et al. (1993) Treffers R. R., et al., 1993, International Astronomical Union Circular, 5870, 3
  • Treffers et al. (1994) Treffers R. R., et al., 1994, International Astronomical Union Circular, 5946, 2
  • Tsvetkov (1982) Tsvetkov D. Y., 1982, Soviet Astronomy Letters, 8, 115
  • Tsvetkov (1986) Tsvetkov D. Y., 1986, Peremennye Zvezdy, 22, 279
  • Tsvetkov & Pavlyuk (1997) Tsvetkov D. Y., et al., 1997, Astronomy Letters, 23, 26
  • Tsvetkov et al. (1990) Tsvetkov D. Y., et al., 1990, A&A, 236, 133
  • Tsvetkov et al. (2013) Tsvetkov D. Y., et al., 2013, Contributions of the Astronomical Observatory Skalnate Pleso, 43, 94
  • Tsvetkov et al. (2014) Tsvetkov D. Y., et al., 2014, Contributions of the Astronomical Observatory Skalnate Pleso, 44, 67
  • Tucker et al. (2018) Tucker M. A., et al., 2018, arXiv e-prints, p. arXiv:1811.09635
  • Tully et al. (2008) Tully R. B., et al., 2008, ApJ, 676, 184
  • Turatto et al. (1990) Turatto M., et al., 1990, AJ, 100, 771
  • Turatto et al. (1996) Turatto M., et al., 1996, MNRAS, 283, 1
  • Tutukov & Yungelson (1979) Tutukov A. V., et al., 1979, Acta Astron., 29, 665
  • Uddin et al. (2017) Uddin S., et al., 2017, The Astronomer’s Telegram, 10605, 1
  • Valenti et al. (2017) Valenti S., et al., 2017, Transient Name Server Classification Report, 2017-613, 1
  • Valentini et al. (2003) Valentini G., et al., 2003, ApJ, 595, 779
  • Vallely et al. (2019) Vallely P., et al., 2019, Bimodal SNe Ia in the Nebular Phase
  • Vanmunster et al. (1994) Vanmunster T., et al., 1994, International Astronomical Union Circular, 6115, 3
  • Verheijen & Sancisi (2001) Verheijen M. A. W., et al., 2001, A&A, 370, 765
  • Vernet et al. (2011) Vernet J., et al., 2011, A&A, 536
  • Vettolani et al. (1981) Vettolani G., et al., 1981, International Astronomical Union Circular, 3584, 1
  • Villi et al. (1998) Villi M., et al., 1998, International Astronomical Union Circular, 6899, 1
  • Villi et al. (2008) Villi M., et al., 2008, Central Bureau Electronic Telegrams, 1228, 1
  • Vinkó et al. (2003) Vinkó J., et al., 2003, A&A, 397, 115
  • Vladimirov et al. (2015) Vladimirov V., et al., 2015, The Astronomer’s Telegram, 7732, 1
  • Volkov (1991) Volkov I. M., 1991, Information Bulletin on Variable Stars, 3581, 1
  • Waagen et al. (1991) Waagen E., et al., 1991, International Astronomical Union Circular, 5239, 1
  • Walker et al. (1994) Walker A., et al., 1994, International Astronomical Union Circular, 5950, 1
  • Walker et al. (2015) Walker E. S., et al., 2015, The Astrophysical Journal Supplement Series, 219
  • Wang et al. (2004) Wang L., et al., 2004, ApJ, 604, L53
  • Wang et al. (2008) Wang X., et al., 2008, ApJ, 675, 626
  • Webbink (1984) Webbink R. F., 1984, ApJ, 277, 355
  • Wells et al. (1994) Wells L. A., et al., 1994, AJ, 108, 2233
  • Weyant et al. (2018) Weyant A., et al., 2018, AJ, 155
  • Wheeler et al. (1975) Wheeler J. C., et al., 1975, ApJ, 200, 145
  • Whelan & Iben (1973) Whelan J., et al., 1973, ApJ, 186, 1007
  • Wiethoff et al. (2017) Wiethoff W., et al., 2017, The Astronomer’s Telegram, 10521, 1
  • Wood-Vasey et al. (2002a) Wood-Vasey W. M., et al., 2002a, International Astronomical Union Circular, 7902, 3
  • Wood-Vasey et al. (2002b) Wood-Vasey W. M., et al., 2002b, International Astronomical Union Circular, 7959, 1
  • Woods et al. (2006) Woods D. F., et al., 2006, AJ, 132, 197
  • Woods et al. (2018) Woods T. E., et al., 2018, ApJ, 863, 120
  • Woosley & Weaver (1994) Woosley S. E., et al., 1994, ApJ, 423, 371
  • Wyrzykowski et al. (2015) Wyrzykowski L., et al., 2015, The Astronomer’s Telegram, 8484, 1
  • Yamanaka et al. (2015) Yamanaka M., et al., 2015, ApJ, 806, 191
  • Yaron & Gal-Yam (2012) Yaron O., et al., 2012, Publications of the Astronomical Society of the Pacific, 124, 668
  • Yoon & Langer (2003) Yoon S. C., et al., 2003, A&A, 412, L53
  • Yu et al. (2000) Yu C., et al., 2000, International Astronomical Union Circular, 7458, 1
  • Zhai et al. (2016) Zhai Q., et al., 2016, AJ, 151
  • Zhang & Wang (2014) Zhang J., et al., 2014, The Astronomer’s Telegram, 6813, 1
  • Zhang et al. (2011) Zhang T., et al., 2011, Central Bureau Electronic Telegrams, 2708, 3
  • Zhang et al. (2014) Zhang J.-J., et al., 2014, AJ, 148
  • Zhang et al. (2018) Zhang J., et al., 2018, Transient Name Server Classification Report, 2018-293, 1
  • Zheng et al. (2013) Zheng W., et al., 2013, The Astronomer’s Telegram, 5637, 1
  • de Vaucouleurs et al. (1991) de Vaucouleurs G., et al., 1991, Third Reference Catalogue of Bright Galaxies
  • van Driel et al. (2001) van Driel W., et al., 2001, A&A, 378, 370
  • van Dyk et al. (1994) van Dyk S. D., et al., 1994, International Astronomical Union Circular, 6105, 1
  • van Genderen (1975) van Genderen A. M., 1975, A&A, 45, 429

Appendix A A Nebular Phase Phillip’s Relation

Figure 8: Left: Residuals of the late-time relation bootstrap fit from Eqs. 6-8 using the values in Table 4. Right: Collapsed residual distribution of the best-fit solution.
Filter
[ mag day] [mag]
42 346 0.00 0.08
67 438 0.03 0.13
34 286 0.00 0.09
Table 4: Values of the coefficients for Eqs. 6-8 and fit statistics. is the total number of photometric points used in each filter from SNe Ia.

In most SNe Ia, the peak luminosity and photospheric phase decline rate (e.g., ) are correlated with the amount of Ni produced in the explosion (e.g., Stritzinger et al., 2006; Scalzo et al., 2019). Therefore, these same observables should also correlate with the magnitude of SNe Ia as they enter the nebular phase. For SNe Ia with nebular spectra but no usable nebular photometry, this relation provides a method for estimating the nebular magnitude using near-peak photometry.

The photometric sample used in deriving the NPPR excludes Iax, CSM, and SC SNe Ia. Although 91T- and 91bg-like do not strictly follow the relation between luminosity and decline rate of normal SNe Ia, they are powered by the radioactive decay of Ni to stable Fe. As mentioned before, is indicative of the amount of Ni synthesised in the explosion, and therefore our parameterization described below still accurately models 91T- and 91bg-like SNe Ia. However, SC and Iax SNe Ia have unique ionisation properties, which is further exemplified by their nebular spectra which lack the prominent [FeII/III] and [CoII/III] emission features of their normal, 91T-, and 91bg-like counterparts (e.g., Taubenberger et al., 2013a; Foley et al., 2016). It is possible that the photometric intricacies of 91T- and 91bg-like SNe Ia are washed out by our heterogeneous sample, and more precise results can be attained with distinct samples of SNe Ia spectral types.

Taking all available nebular phase photometry of viable SNe Ia from this work and the literature, we derive an approximate functional form for calculating the apparent magnitude of a SN Ia with a measured and . Since SNe Ia have nearly linear decays in magnitude space at nebular epochs, we use the functional form

(5)

where is the nebular magnitude in filter at phase , the magnitude at peak in filter , and is the change in brightness between maximum light and for that filter. By formulating our relation for individual filters, we can neglect extinction from the Milky Way and the host galaxy since the maximum light and nebular magnitude of a SN Ia will be affected equally. Thus, we parameterize as a function a SN Ia’s ,

(6)

where

(7)

and

(8)

Here, is the apparent magnitude at maximum light in filter , is the nebular magnitude, is the phase of the observations, the decline rate, and measured coefficients which are provided in Table 4. and are offset by typical values to reduce their covariance in the fitting process.

The coefficients in Table 4 were computed using all available nebular photometry between after maximum light. The coefficients were first approximated using a sample of well-studied SNe Ia with measurements in a given filter in the temporal bounds listed above, such as SNe 2011fe, 2012fr, 2013gy, and 2015F, then expanded to include all photometric points. The SNe Ia used in deriving the NPPR have decline rates that span and are denoted with a in Table 16. For publicly available photometry for which there are no reported uncertainties, we assign a nominal uncertainty of . In fitting the data, we implement non-linear least squares fitting coupled with a bootstrap-resampling technique to derive reasonable estimates for the uncertainties. The residuals of the best-fit solution are shown in the left panel of Fig. 8, and the collapsed distribution is provided in the right panel.

For SNe Ia with a measured peak magnitude and , we show the nebular magnitude can be approximated to . These results were derived using a heterogeneous data set and likely can be improved with a consistent photometric system and targeted observations across a reasonable span of . This technique can also be used in identifying peculiar or strange SNe Ia that deviate from their expected brightness at a given epoch, such as “late-onset” CSM interaction (e.g., Graham et al., 2019). Additionally, we attempted to expand this methodology to other photometric filters (e.g., , ), but there were too few observations to build a quality model.

Appendix B Supplementary Tables and Figures

In Table 6, we provide the name of the SN Ia, redshift, and references for discovery and classification. For information on each spectrum, including the date, telescope, instrument, and reference, see Table 8. Flux limits and derived mass limits are given in Table 12. New photometry presented in this work is provided in Table 14 and all photometry references are given in Table 16.

b.1 Data Tables

For SNe Ia with redshifts measured from the supernova lines near maximum light, we tweak the redshift using host galaxy emission lines when necessary. Major host galaxy lines such as H, H, [NII], and [OIII] are fit with Gaussian line profiles to estimate the line centre and then used to measure the host redshift.

For SNe Ia with insufficient photometry for a reliable light curve fit in SNooPy, we consider two approaches. If there are photometric points near maximum light, we use linear least-squares coupled with bootstrap-resampling to fit a quadratic curve to the data and estimate and the associated uncertainty. Otherwise, the value for is taken from the spectroscopic classification reference given in Table 6 and assigned a nominal uncertainty of .

SNe-Iax do not conform to the standard SNe Ia templates utilised by SNooPy and other SN Ia light curve fitters. Thus, we compute and using spline fits in the SNooPy environment. This prevents us from deriving the host reddening , however, the Iax SNe in our sample have negligible reddening (Li et al., 2003a; Phillips et al., 2007; McCully et al., 2014; Stritzinger et al., 2015; Foley et al., 2015).

The “Quality” column in Table 8 provides a rough estimate of the quality of the spectrum. This is mostly qualitative, and intended to provide readers with an estimate of the spectral quality for each SN Ia in our sample. The rankings are as follows:

  • High: The spectrum clearly shows the major Fe and Co emission lines between Å. The spectrum exhibits little to no host contamination or instrumental artefacts.

  • Med(ium): The major Fe lines are visible, while the Co lines are noisy or absent. The spectrum may also suffer from minor to moderate host galaxy contamination and/or instrumental artefacts.

  • Low: The major Fe lines are barely detectable above the spectral noise, and the Co lines mostly below the detection threshold. This category also includes overall medium-quality spectra with significant host galaxy contamination and/or instrumental artefacts.

b.2 Special Cases

Name Ref.
SN2002bo 1.2 Phillips et al. (2013)
SN2004eo 0.8 Burns et al. (2014)
SN2006X 1.5 Wang et al. (2008); Phillips et al. (2013); Burns et al. (2014)
SN2007le 1.6 Phillips et al. (2013); Burns et al. (2014)
SN2014J 1.5 Amanullah et al. (2014); Foley et al. (2014); Gao et al. (2015); Brown et al. (2015)
Table 5: values and references for SNe Ia with deviations from the assumed .

We discuss any extenuating circumstances or any other relevant details about specific SNe Ia that differ from the general methodology described in §2. Examples include alternative flux calibration methods, spectroscopic oddities noticed in our analysis, and spectrum reference discrepancies. SNe Ia with values known to deviate from the standard are listed in Table 5. For ensemble studies (e.g., Phillips et al., 2013; Burns et al., 2014), we require a deviation from to include the value in our calculation. When drawing values from Burns et al. (2014), we implement the F99+uniform prior results.

SN1981B: The nebular spectrum on the OSC and WISeREP refer to Branch et al. (1983) as the source of the spectrum, however, we find no mention therein. Therefore, we include this reference in Table 8, with the caveat that we could not verify its source. Additionally, there are no definitive narrow features in the spectrum to confidently determine the spectral resolution. We assume a resolution of Å across the entire spectral range.

SN1998bu: The two nebular spectra from Cappellaro et al. (2001) do not have any specific mention in that manuscript, however, the reference on WiseRep points to this paper. Thus, we include the reference, but acknowledge we could not verify this paper was the true source for these spectra.

SN2002bo: The OSC and WiseRep also report several nebular phase NIR spectra for this SN. However, cross-referencing the reported spectra with the observational parameters given in Benetti et al. (2004), we believe the dates provided for the NIR spectra are off by a year, and these spectra are closer to a few months after maximum light instead of several hundred days after maximum light. We exclude these spectra from our sample.

b.3 Supplementary Figures

We provide cutouts around each spectral line inspected for H/He emission (Table 2) for the spectrum used in calculating the limits provided in Table 12 for each SN Ia as supplementary figures. An example of the format of these figures is provided in Fig. 3. The black line is the observed spectrum, with the continuum fit and flux upper limit in red and purple, respectively. Gray shaded areas indicate masked spectral regions and completely gray boxes indicate that particular SN Ia had no spectra covering that spectral region. When multiple nebular spectra of a SN Ia cover the same expected H/He line, we provide the spectrum corresponding to the best mass limit for that line. Therefore, the panel for H may show a different spectrum than the panel for H for the same SN Ia. This ensures all adequate spectra are presented, even when some spectra do not cover all the optical and NIR lines considered in this study. The border colour of a given panel indicates whether it is used in the final stripped mass determination, the results of which are provided in Table 12. Blue borders indicate the panels used in the H-rich mass limit, red borders indicate He-rich limits, and purple borders indicate He lines used for both H- and He-rich limits.

Disc. Name IAU Name Pec? z Redshift Ref. Discovery Classification
ASASSN-14hr N Jones et al. (2009) Nicolas et al. (2014) Morrell et al. (2014a)
ASASSN-14jc N Jones et al. (2009) Kiyota et al. (2014a) Romero-Canizales et al. (2014)
ASASSN-14jg N Jones et al. (2009) Holoien et al. (2014) Arcavi et al. (2014)
ASASSN-14jz N This Work Kiyota et al. (2014b) Dimitriadis et al. (2014)
ASASSN-14kq N Jones et al. (2009) Brimacombe et al. (2014a) Morrell et al. (2014b)
ASASSN-14lt N Springob et al. (2014) Kiyota et al. (2014c) Zhang & Wang (2014)
ASASSN-14lu N Collobert et al. (2006) Brimacombe et al. (2014b) Shappee et al. (2014b)
ASASSN-14lv N Jones et al. (2006) Kiyota et al. (2014d) Shappee et al. (2015)
ASASSN-14me N Shappee et al. (2015) Stanek et al. (2014) Shappee et al. (2015)
ASASSN-15be N Colless et al. (2003) Brimacombe et al. (2015) Morrell et al. (2015)
ASASSN-15hx N This Work Dong et al. (2015b) Frohmaier et al. (2015)
PSN J1149 N Meyer et al. (2004) Vladimirov et al. (2015) Rudy et al. (2015)
SN1972E N Koribalski et al. (2004) Kowal (1972) Herbig et al. (1972)
SN1981B N Grogin et al. (1998) Aksenov (1981) Vettolani et al. (1981)
SN1986G 91bg-like Graham et al. (1978) Evans et al. (1986) Feast et al. (1986)
SN1990N N Meyer et al. (2004) Maury et al. (1990) Maury et al. (1990)
SN1991T 91T-like Strauss et al. (1992) Waagen et al. (1991) Waagen et al. (1991)
SN1991bg 91bg-like Cappellari et al. (2011) Kosai et al. (1991) Kirshner et al. (1991)
SN1992A N D’Onofrio et al. (1995) Liller et al. (1992) Liller et al. (1992)
SN1993Z N Epinat et al. (2008) Treffers et al. (1993) Treffers et al. (1993)
SN1994ae N Krumm & Salpeter (1980) van Dyk et al. (1994) Iijima et al. (1994)
SN1995D N Cappellari et al. (2011) Nakano et al. (1995) Benetti et al. (1995)
SN1996X N Ogando et al. (2008) Garradd et al. (1996) Suntzeff et al. (1996)
SN1998aq N de Vaucouleurs et al. (1991) Hurst et al. (1998) Ayani & Yamaoka (1998)
SN1998bu N de Vaucouleurs et al. (1991) Villi et al. (1998) Ayani et al. (1998)
SN1999aa 91T-like de Vaucouleurs et al. (1991) Armstrong & Schwartz (1999) Filippenko et al. (1999)
SN1999by 91bg-like de Vaucouleurs et al. (1991) Arbour et al. (1999) Gerardy & Fesen (1999)
SN2000cx 91T-like Cappellari et al. (2011) Yu et al. (2000) Chornock et al. (2000)
SN2001el N Koribalski et al. (2004) Monard et al. (2001) Sollerman et al. (2001)
SN2002bo N Theureau et al. (1998) Cacella et al. (2002) Kawakita et al. (2002)
SN2002cx Iax Falco et al. (1999) Wood-Vasey et al. (2002a) Matheson et al. (2002)
SN2002dj N Rothberg & Joseph (2006) Hutchings & Li (2002) Riello et al. (2002)
SN2002er N de Vaucouleurs et al. (1991) Wood-Vasey et al. (2002b) Smartt et al. (2002)
SN2003cg N van Driel et al. (2001) Nakano et al. (2003) Kotak et al. (2003a)
SN2003du N Schneider et al. (1992) Schwartz & Holvorcem (2003) Kotak et al. (2003b)
SN2003gs 91bg-like Smith et al. (2000) Evans et al. (2003) Evans et al. (2003)
SN2003hv N Ogando et al. (2008) Beutler & Li (2003) Dressler et al. (2003)
SN2003kf N Theureau et al. (1998) Li et al. (2003b) Li et al. (2003b)
SN2004S N Theureau et al. (2007) Martin & Biggs (2004) Suntzeff et al. (2004)
SN2004eo N Theureau et al. (1998) Arbour et al. (2004) Gonzalez et al. (2004)
SN2005W N de Vaucouleurs et al. (1991) Nakano et al. (2005) Elias-Rosa et al. (2005)
SN2005am N Theureau et al. (1998) Martin et al. (2005) Modjaz et al. (2005b)
SN2005cf N de Vaucouleurs et al. (1991) Pugh & Li (2005) Modjaz et al. (2005a)
SN2005hk Iax Abolfathi et al. (2018) Quimby et al. (2005) Serduke et al. (2005)
SN2006X N Rand (1995) Ponticello et al. (2006) Quimby et al. (2006a)
SN2006dd N Longhetti et al. (1998) Quimby et al. (2006b) Salvo et al. (2006)
SN2006gz SC Falco et al. (1999) Frieman et al. (2006) Prieto et al. (2006)
SN2007af N Koribalski et al. (2004) Nakano & Itagaki (2007) Salgado et al. (2007)
SN2007if SC Scalzo et al. (2010) Akerlof et al. (2007) Akerlof et al. (2007)
SN2007le N Koribalski et al. (2004) Monard et al. (2007) Filippenko et al. (2007)
SN2007on N Graham et al. (1998) Pollas & Klotz (2007) Morrell et al. (2007)
SN2008A Iax Theureau et al. (1998) Nakano et al. (2008) Blondin & Berlind (2008)
SN2008Q N Cappellari et al. (2011) Villi et al. (2008) Stanishev & Pursimo (2008)
SN2009dc SC Falco et al. (1999) Puckett et al. (2009) Harutyunyan et al. (2009)
SN2009ig N Meyer et al. (2004) Kleiser et al. (2009) Kleiser et al. (2009)
SN2009le N Theureau et al. (1998) Pignata et al. (2009) Challis & Berlind (2009)
SN2010ev N Meyer et al. (2004) Pignata et al. (2010) Stritzinger (2010)
SN2010gp N Downes et al. (1993) Maza et al. (2010) Folatelli et al. (2010b)
SN2010hg N Meyer et al. (2004) Monard & Africa (2010) Prieto et al. (2010)
SN2010lp 91bg-like Huchra et al. (1999) Cox et al. (2010) Prieto & Morrell (2011)
SN2011K N Monard et al. (2011) Drake et al. (2011a) Drake et al. (2011a)
Table 6: All SNe Ia studied in this work.

All SNe Ia studied in this work. Disc. Name IAU Name Pec? z Redshift Ref. Discovery Classification SNhunt37 SN2011ae N Meyer et al. (2004) Howerton et al. (2011) Howerton et al. (2011) SN2011at N Theureau et al. (1998) Cox et al. (2011) Cox et al. (2011) SN2011by N Verheijen & Sancisi (2001) Drake et al. (2011b) Zhang et al. (2011) SN2011ek N Rhee & van Albada (1996) Nakano et al. (2011) Nakano et al. (2011) PTF11kly SN2011fe N Maguire et al. (2014) Nugent et al. (2011b) Cenko et al. (2011) SN2011im N Catinella et al. (2005) Brimacombe et al. (2011) Brimacombe et al. (2011) SN2011iv N Graham et al. (1998) Noguchi et al. (2011a) Noguchi et al. (2011a) SN2011iy N Corsini et al. (2003) Noguchi et al. (2011b) Noguchi et al. (2011b) SN2012Z Iax Koribalski et al. (2004) Cenko et al. (2012) Cenko et al. (2012) SN2012cg N Kent et al. (2008) Kandrashoff et al. (2012) Kandrashoff et al. (2012) SNhunt136 SN2012cu N de Vaucouleurs et al. (1991) Marion et al. (2012) Marion et al. (2012) SN2012ei N Galbany et al. (2014) Nakano et al. (2012) Nakano et al. (2012) SN2012fr N Bureau et al. (1996) Klotz et al. (2012) Klotz et al. (2012) SN2012hr N Tully et al. (2008) Drescher et al. (2012) Drescher et al. (2012) SN2012ht N Guthrie & Napier (1996) Nishiyama et al. (2012) Nishiyama et al. (2012) SN2013aa N Huchra et al. (1999) Parker et al. (2013a) Parker et al. (2013a) SNhunt196 SN2013cs N Pisano et al. (2011) Howerton et al. (2013) Howerton et al. (2013) SN2013ct N Smoker et al. (2000) Parker et al. (2013b) Parker et al. (2013b) SN2013dy N Pan et al. (2015) Casper et al. (2013) Casper et al. (2013) SN2013gy N Catinella et al. (2005) Kim et al. (2013) Kim et al. (2013) SN2014J N de Vaucouleurs et al. (1991) Fossey et al. (2014) Itagaki et al. (2014) SN2014bv N de Vaucouleurs et al. (1991) Cortini et al. (2014) Cortini et al. (2014) SN2015F N Meyer et al. (2004) Monard et al. (2015) Monard et al. (2015) SN2015I N Giovanelli et al. (1997) Nakano et al. (2015) Karamehmetoglu et al. (2015) SN2016brx 91bg-like Jones et al. (2009) Parker (2016) Morrell & Shappee (2016) SN2016bry N Rhee & van Albada (1996) Tinella (2016) Ochner et al. (2016) ASASSN-16eq SN2016bsa N Paturel et al. (2003) Brimacombe et al. (2016) Mattila et al. (2016) Gaia16avm SN2016ehy N Halevi et al. (2016) Delgado et al. (2016) Halevi et al. (2016) ATLAS16cpu SN2016ffh N Abazajian et al. (2005) Tonry et al. (2016) Terreran et al. (2016) SN2016gxp 91T-like Huchra et al. (1999) Gagliano et al. (2016) Leadbeater (2016) ASASSN-17cs SN2017azw N Nyholm et al. (2017) Brimacombe et al. (2017) Nyholm et al. (2017) DLT17u SN2017cbv N Koribalski et al. (2004) Tartaglia et al. (2017a) Hosseinzadeh et al. (2017a) ATLAS17dfo SN2017ckq N Mathewson et al. (1992) Tonry et al. (2017c) Pan et al. (2017) DLT17ar SN2017cyy N Meyer et al. (2004) Tartaglia et al. (2017b) Jha et al. (2017) DLT17bk SN2017ejb 91bg-like de Vaucouleurs et al. (1991) Tartaglia et al. (2017c) Valenti et al. (2017) ASASSN-18hz SN2017evn N Adelman-McCarthy et al. (2008) Wiethoff et al. (2017) Everson et al. (2017) SN2017ezd N Jones et al. (2009) Parker (2017) Uddin et al. (2017) DLT17bx SN2017fgc N Cappellari et al. (2011) Sand et al. (2017) Sand et al. (2017) DLT17cd SN2017fzw 91bg-like de Vaucouleurs et al. (1991) Tartaglia et al. (2017d) Hosseinzadeh et al. (2017b) ATLAS17jiv SN2017gah N Lauberts & Valentijn (1989) Tonry et al. (2017d) Hosseinzadeh et al. (2017b) SN2017glq N Woods et al. (2006) Gagliano et al. (2017) Prentice & Ashall (2017) ATLAS17nmh SN2017isq N Benetti et al. (2017) Tonry et al. (2017a) Benetti et al. (2017) ATLAS17nse SN2017iyb N Meyer et al. (2004) Tonry et al. (2017b) Jones et al. (2017) ASASSN-18hb SN2018aqi N Theureau et al. (1998) Malesani et al. (2018) Brimacombe et al. (2018) ASASSN-18bt SN2018oh N Schneider et al. (1990) Brown et al. (2018) Leadbeater (2018) ASASSN-18da SN2018vw 91T-like Dong et al. (2018b) Stanek (2018) Stone et al. (2018) DLT18h SN2018xx N Smith et al. (2000) Sand et al. (2018b) Sand et al. (2018c) DLT18i SN2018yu N de Vaucouleurs et al. (1991) Sand et al. (2018d) Zhang et al. (2018) SNF-012 SC Taubenberger et al. (2013a)

  • PSN J1149 = PSN J11492548-0507138

  • SNF-012 = SNF20080723-012

Table 7: continued
SN Obs. Date Phase Telescope Instrument Range Expt. Qual. Ref.
(MJD) (days) (Å) (s) [km s]
SN1972E 41651.00 205.1 P200 DBSP Med Kirshner & Oke (1975)
41682.00 236.1 P200 DBSP Med Kirshner & Oke (1975)
41795.00 349.1 P200 DBSP Med Kirshner & Oke (1975)
41864.00 418.1 P200 DBSP Med Kirshner & Oke (1975)
SN1981B 44940.00 267.3 Med Branch et al. (1993)
SN1986G 46816.50 255.5 ESO2.2m BC Med Cristiani et al. (1992)
46817.50 256.5 ESO3.6m EFOSC High Cristiani et al. (1992)
46818.50 257.5 ESO3.6m EFOSC High Cristiani et al. (1992)
46885.50 324.5 ESO3.6m EFOSC Med Cristiani et al. (1992)
SN1990N 48268.50 186.0 WHT FOS High Gomez et al. (1996)
48309.50 227.0 WHT FOS High Gomez et al. (1996)
48337.50 255.0 WHT FOS Med Gomez et al. (1996)
48362.50 280.0 WHT FOS Med Gomez et al. (1996)
48415.50 333.0 WHT FOS Med Gomez et al. (1996)
SN1991T 48632.50 258.0 WHT ISIS High Gomez et al. (1996)
48658.50 284.0 INT FOS1 Med Gomez et al. (1996)
48690.50 316.0 INT FOS1 Med Gomez et al. (1996)
48694.93 320.4 Shane3m KAST Med BSNIP
48723.86 349.3 Shane3m KAST Med BSNIP
SN1991bg 48802.50 198.4 WHT FOS2 Low Gomez et al. (1996)
48806.50 202.4 ESO3.6m EFOSC2 Med Turatto et al. (1996)
SN1992A 48931.00 291.7 HST FOS Low Kirshner et al. (1993)
SN1993Z 49428.50 181.4 Shane3m KAST Med BSNIP
49460.00 212.9 Shane3m KAST Med BSNIP
SN1994ae 49904.71 219.4 Shane3m KAST Med BSNIP
50053.97 368.7 MMT BCS Med CfA
SN1995D 50045.00 277.6 MMT BCS Med CfA
50053.00 285.6 MMT BCS Med CfA
SN1996X 50437.00 247.2 ESO1.5m BC Low Salvo et al. (2001)
50489.00 299.2 ESO3.6m EFOSC2 High Salvo et al. (2001)
SN1998aq 51141.00 210.2 Till FAST Med CfA
51161.00 230.2 Till FAST Med CfA
51171.00 240.2 Till FAST Med CfA
SN1998bu 51143.04 191.2 Till FAST High CfA
51161.04 209.2 Till FAST High CfA
51169.98 218.1 Till FAST Med CfA
51188.88 237.0 Shane3m KAST High BSNIP
51195.96 244.1 Till FAST Med CfA
51201.96 250.1 D1.54m DFOSC Med Cappellaro et al. (2001)
51204.96 253.1 NTT SOFI Med Spyromilio et al. (2004)
51232.92 281.0 Shane3m KAST High BSNIP
51281.50 329.6 ESO3.6m EFOSC2 Med Cappellaro et al. (2001)
51292.78 340.9 Shane3m KAST Med BSNIP
SN1999aa 51491.00 258.5 KeckI LRIS Med BSNIP
51517.00 284.5 KeckI LRIS Med BSNIP
SN1999by 51494.14 185.2 KeckI LRIS High BSNIP
SN2000cx 51935.00 182.2 MMT BCS Med CfA
52204.00 451.2 KeckII ESI Low BSNIP
SN2001el 52492.40 310.1 VLT FORS1 Med PI-Sollerman
52500.35 318.1 VLT FORS1 Med PI-Sollerman
52580.15 397.9 VLT FORS1 Med Mattila et al. (2005)
SN2002bo 52584.00 227.2 KeckII ESI low BSNIP
52668.00 311.2 MMT BCS Med Blondin et al. (2012)
SN2002cx 52647.15 232.1 KeckI LRIS Low BSNIP
52699.11 284.0 KeckI LRIS Low BSNIP
52699.14 284.1 KeckI LRIS Low BSNIP
52731.81 316.7 Clay LDSS Low CfA
SN2002dj 52671.00 219.7 NTT EFOSC2 Med Pignata et al. (2008)
52724.00 272.7 VLT FORS1 Med Pignata et al. (2008)
Table 8: Spectra observations.

Spectra observations. SN Obs. Date Phase Telescope Instrument Range Expt. Qual. Ref. (MJD) (days) (Å) (s) [km s] SN2002er 52739.00 213.7 TNG DOLORES