An impostor among us I: Photometric and spectroscopic evolution of AT 2016jbu
S. J. Brennan, M. Fraser, J. Johansson, A. Pastorello, R. Kotak, H. F. Stevance, T. -W. Chen, J. J. Eldridge, S. Bose, P. J. Brown, E. Callis, R. Cartier, M. Dennefeld, Subo Dong, P. Duffy, N. Elias-Rosa, G. Hosseinzadeh, E. Hsiao, H. Kuncarayakti, A. Martin-Carrillo, B. Monard, A. Nyholm, G. Pignata, D. Sand, B. J. Shappee, S. J. Smartt, B. E. Tucker, L. Wyrzykowski, H. Abbot, S. Benetti, S. Blondin, Ping Chen, J. Bento, A. Delgado, L. Galbany, M. Gromadzki, C. P. Gutiérrez, L. Hanlon, D. L. Harrison, D. Hiramatsu, S. T. Hodgkin, T. W. -S. Holoien, D. A. Howell, C. Inserra, E. Kankare, S. Kozlowski, K. Maguire, T. E. Müller-Bravo, C. McCully, P. Meintjes, N. Morrell, M. Nicholl, D. O'Neill, P. Pietrukowicz, R. Poleski, J. L. Prieto, A. Rau, D. E. Reichart, T. Schweyer, M. Shahbandeh, J. Skowron, J. Sollerman, I. Sosz?yski, M. D. Stritzinger, M. Szyma?ski, L. Tartaglia, A. Udalski, K. Ulaczyk, D. R. Young, M. van Leeuwen, B. van Soelen
MMNRAS , 1–28 (2021) Preprint 3 March 2021 Compiled using MNRAS L A TEX style file v3.0
An impostor among us I:Photometric and spectroscopic evolution of AT 2016jbu
S. J. Brennan (cid:63) , M. Fraser , J. Johansson , A. Pastorello , R. Kotak H. F. Stevance ,T. -W. Chen , , J. J. Eldridge , S. Bose , , P. J. Brown , E. Callis , R. Cartier ,M. Dennefeld , Subo Dong , P. Duffy , N. Elias-Rosa , , G. Hosseinzadeh ,E. Hsiao , H. Kuncarayakti , , A. Martin-Carrillo , B. Monard , A. Nyholm , G. Pignata , ,D. Sand , B. J. Shappee , S. J. Smartt , B. E. Tucker , , , L. Wyrzykowski ,H. Abbot , S. Benetti , J. Bento , S. Blondin , , Ping Chen , A. Delgado , , L. Galbany ,M. Gromadzki , C. P. Guti´errez , , L. Hanlon , D. L. Harrison , , D. Hiramatsu , ,S. T. Hodgkin , T. W. -S. Holoien , D. A. Howell , , C. Inserra , E. Kankare ,S. Koz(cid:32)lowski , T. E. M¨uller-Bravo , K. Maguire , C. McCully , , P. Meintjes ,N. Morrell , M. Nicholl , , D. O’Neill , P. Pietrukowicz , R. Poleski , J. L. Prieto , ,A. Rau , D. E. Reichart , T. Schweyer , , M. Shahbandeh , J. Skowron ,J. Sollerman , I. Soszy´nski , M. D. Stritzinger , M. Szyma´nski , L. Tartaglia ,A. Udalski , K. Ulaczyk , , D. R. Young , M. van Leeuwen , B. van Soelen The authors’ affiliations are shown in Appendix A.
ABSTRACT
We present comprehensive, multi-wavelength observations of AT 2016jbu, an interacting transient. High cadencephotometric coverage reveals that AT 2016jbu underwent significant photometric variability followed by two luminousevents, the latter of which reached an absolute magnitude of M V ∼ − . − seen innarrow emission features from a slow moving CSM, and up to 10,000 km s − seen in broad absorption from somehigh velocity material. Similar velocities are seen in other SN 2009ip-like transients. Late-time spectra ( ∼ +1 year)show a lack of forbidden emission lines expected from a core-collapse supernova during the nebular phase and aredominated by strong emission from H, He i and Ca ii . Strong asymmetric emission features, a bumpy lightcurve, andcontinually evolving spectra suggest late time CSM interaction is inhibiting the emergence of any nebular features.We compare the evolution of H α among SN 2009ip-like transients and find possible evidence for orientation angleeffects. The light-curve evolution of AT 2016jbu suggests similar, but not identical, circumstellar environments toother SN 2009ip-like transients. In Paper II we continue the discussion of AT 2016jbu and SN 2009ip-like transientsand using the data presented here, we focus on the local environment, the progenitor, and on modelling the transientitself. Key words: circumstellar matter – stars: massive – supernovae: individual: AT 2016jbu – supernovae: individual:Gaia16cfr – supernovae: individual: SN 2009ip
Massive stars are among the drivers of the dynamic evolu-tion of galaxies. They seed the interstellar medium with the (cid:63)
Contact e-mail: [email protected] products of stellar nucleosynthesis, and inject kinetic energyboth through their winds and when they eventually explodeas supernovae. Despite their importance, the lives and deathsof massive stars remain relatively poorly understood. In par-ticular, it is not known with any degree of certainty whichmassive stars produce which supernovae, at what stage of © a r X i v : . [ a s t r o - ph . S R ] M a r S. J. Brennan et al. their evolution, and what occurs in the ∼ years prior tocore-collapse.Massive stars that eventually undergo core-collapse whensurrounded by some dense circumstellar material (CSM) areknown as Type IIn supernovae (SNe) (Schlegel 1990; Fil-ippenko 1997; Fraser 2020). This is signified in spectra bya bright, blue continuum with narrow H and He i emissionlines at early times. Type IIn SNe spectra show narrow( ∼ −
500 km s − ) components arising in the photo-ionised, slow moving CSM. Intermediate width emission lines( ∼ − ) arise from either electron scattering of pho-tons in narrower lines or emission from gas shocked by super-nova (SN) ejecta. Some events also show very broad emissionor absorption features ( ∼ ,
000 km s − ) arising from fastejecta, typically associated with material ejected in a core-collapse explosion.The existence of the dense CSM indicates that the TypeIIn progenitors have high mass-loss rates shortly before theirterminal explosion. This dense material at the end of a star’slife can come from several pathways (see reviews by Puls et al.2008; Smith 2014; Fraser 2020, for further detail.).Luminous Blue Variables (LBVs) are classically believed tobe massive stars ( M ZAMS (cid:38)
25 M (cid:12) ) undergoing a brief phasein evolution, transitioning between the end of the main se-quence and beginning of He-core burning stage (Conti 1984).A well-known attribute of LBV stars are their giant eruptionswith extreme mass loss and increased luminosity over a shortperiod of time. In the past 15 years, there have been a num-ber of suggestions that LBVs may actually explode directly asSNe, including for SNe 2003bg and 2005gj as well as the TypeIIn SN, SN 2005gl (Kotak & Vink 2006; Gal-Yam et al. 2007;Trundle et al. 2008). This contradicts the traditional view onmassive star evolution, where a LBV has just begun He-coreburning, and does not have a massive Fe core necessary toexplode as a CCSN. Further controversy has arisen over theevolutionary status of LBVs, with Smith & Tombleson (2015)suggesting that they are a product of binary evolution in aneffort to explain an apparent discrepancy between the spa-tial distribution of LBVs and their antecedents. However, thestatistical significance of this purported discrepancy has sincebeen disputed by Davidson et al. (2016).Further complicating this picture are a growing numberof extragalactic transients that show narrow emission linesin their spectra (indicating CSM) but have much fainter ab-solute magnitudes than most typical Type IIn SNe. Theseevents are often termed
SN Impostors (Van Dyk et al. 2000;Maund et al. 2006; Pastorello & Fraser 2019), and are be-lieved in many cases to be extra-galactic LBVs experiencinggiant eruptions (e.g. SN 2000ch; Wagner et al. 2004; Pas-torello et al. 2010). These eruptions do not completely de-stroy their progenitors.Perhaps the best studied exemplar of the confusion be-tween LBVs, SN impostors, and genuine Type IIn SNe isSN 2009ip. SN 2009ip was found on 2009 August 26 at ∼ (cid:12) (Smith et al. 2010; Foleyet al. 2011). There is much debate on the fate of SN 2009ip.Some argue that SN 2009ip has finally exploded as a genuineType IIn SN during the 2012 outburst (Prieto et al. 2013;Mauerhan et al. 2013). However, other authors remain agnos-tic as to SN 2009ip’s fate as a CCSN, pointing to the absenceof any evidence for nucleosynthesised material in late-timespectra, as well as SN 2009ip not fading significantly belowthe progenitor magnitude (Fraser et al. 2013; Margutti et al.2014; Fraser et al. 2015).Since the discovery of SN 2009ip, a number of remark-ably similar transients have been found. The growing familyof SN 2009ip-like transients share similar spectral and photo-metric evolution. SN 2009ip-like transients have the followingobservable traits. History of variability lasting (at least) ∼
10 years withoutbursts reaching M r ∼ − ± Two bright luminous events with the first peak reachinga magnitude of M r ∼ − ± M r ∼ − ± Spectroscopically similar to a Type IIn SN i.e. narrowemission features and a blue continuum at early times. Restrictive upper limits to the mass of any explosivelysynthesised Ni.In this paper we focus on one such SN 2009ip-like tran-sient. AT 2016jbu (also known as Gaia16cfr; Bose et al. 2017)was discovered at RA. = 07:36:25.96, DEC. = − Gaia satellite on 2016 December 1 with a mag-nitude of G =19.63 (corresponding to an absolute magnitudeof − .
97 mag for our adopted distance modulus). The PublicESO Spectroscopic Survey for Transient Objects (PESSTO)collaboration (Smartt et al. 2015) classified AT 2016jbu asan SN 2009ip-like transient due to its spectral appearanceand apparent slow rise (Fraser et al. 2017). Fraser et al.(2017) also finds that the progenitor of AT 2016jbu seen inarchival Hubble Space Telescope (
HST ) images is consistentwith a massive ( <
30 M (cid:12) ) progenitor. The transient wasindependently discovered by B. Monard in late Decemberwho reported the likely association of AT 2016jbu to its host,NGC 2442. AT 2016jbu is situated to the south of NGC 2442,a spiral galaxy commonly referred to as the
Meathook Galaxy .NGC 2442 has hosted two other SNe including SN 1999ga,a low luminosity Type II SN (Pastorello et al. 2009) andSN 2015F, a Type Ia SN (Cartier et al. 2017). We mark theirrespective locations in Fig. 1. Bose et al. (2017) and Prenticeet al. (2018) reported initial spectroscopic observations andclassification of AT 2016jbu.AT 2016jbu has been studied by Kilpatrick et al. (2018)(hereafter referred to as K18). K18 finds that AT 2016jbuappears similar to a Type IIn SN, with narrow emissionlines and a blue continuum. The
Gaia lightcurve shows thatAT 2016jbu has a double-peaked lightcurve showing two dis-tinct events (we refer to these events as
Event A and
EventB ). This is common in SN 2009ip-like transient with
Event B reaching an absolute magnitude of r ∼ −
18 mag. H α displaysa double-peaked profile a few weeks after maximum bright- MNRAS , 1–28 (2021) hotometric and spectroscopic evolution of AT 2016jbu h m s m s s -69°30'32'34'36' RA D e c h m s m s s -69°30'32'34'36' RA D e c h m s m s s -69°30'32'34'36' RA D e c Figure 1.
Finder chart for AT 2016jbu. Image is a 60s r-band exposure taken with the LCO 1-m. AT 2016jbu is situated to thesouth-east of the spiral galaxy NGC 2442 nucleus and is indicatedwith a red cross reticle in the center of the image. This locationlies on the outskirts of a
Superbubble (Pancoast et al. 2010), with ahigh star formation rate. We also include the location of the TypeIa SN 2015F (blue circle, north west of image center; Cartier et al.2017 ) and the Type II SN 1999ga (green square, south west ofimage center; Pastorello et al. 2009). ness, indicating a complex CSM environment. K18 model H α using a multi-component line profile including a shifted blueemission feature that grows with time, with their final profilesimilar to that of the Type IIn SN 2015bh (Elias-Rosa et al.2016; Th¨one et al. 2017) at late times. Using HST images,spanning 10 years prior to the 2016 transient, K18 reportsthat AT 2016jbu underwent a series of outbursts in the decadeprior, similar to SN 2009ip. K18 report pre-outburst imagesfrom
HST that are consistent with a ∼
18 M (cid:12) progenitor star,with strong evidence that this progenitor is significantly red-dened by circumstellar (CS) dust, which would allow for ahigher mass. Performing dust modelling using Spitzer pho-tometry, K18 find the spectral energy distribution (SED) ∼ . ± .
06 mag, which is a weighted average of the values deter-mined from
HST observations of Cepheids ( µ = 31 . ± . mag ; Riess et al. 2016) and from the SN Ia 2015F( µ = 31 . ± .
14 mag; Cartier et al. 2017). This correspondsto a metric distance of 20 . ± .
58 Mpc. We adopt a redshiftof z=0.00489 from H I Parkes All Sky Survey (Wong et al.2006). The foreground extinction towards NGC 2442 is takento be A V = 0 .
556 mag, from Schlafly & Finkbeiner (2011) via the NASA Extragalactic Database (NED; ). We do notcorrect for any possible host galaxy or circumstellar extinc-tion, however we note that the blue colors seen in the spectraof AT 2016jbu do not point towards significant reddening bydust. We take the V -band maximum during the second, moreluminous event in the lightcurve (as determined through apolynomial fit) as our reference epoch (MJD 57784 . ± . Optical imaging of AT 2016jbu in
BVRri filters was obtainedwith the 3.58m ESO New Technology Telescope (NTT) +EFOSC2, as part of the ePESSTO survey. All images werereduced in the standard fashion using the PESSTO pipeline(Smartt et al. 2015); in brief images were bias and overscansubtracted, flat fielded, before being cleaned of cosmic raysusing a Laplacian filter (van Dokkum 2001). Further opti-cal imaging was obtained from the Las Cumbres Observatorynetwork of robotic 1-m telescopes as part of the Global Super-nova Project. These data were reduced automatically by the banzai pipeline, which runs on all Las Cumbres Observatory(LCO) Global Telescope images (Brown et al. 2013). Imageswere also obtained from the Watcher telescope. Watcher isa 40 cm robotic telescope located at Boyden Observatory inSouth Africa (French et al. 2004). It is equipped with an An-dor IXon EMCCD camera providing a field of view of 8 × Python .AT 2016jbu was monitored using the Gamma-Ray BurstOptical/Near-Infrared Detector (GROND; Greiner et al.(2008)), a 7-channel imager that collects multi-color pho-tometry simultaneously with Sloan- griz and
JHK/Ks bands,mounted at the 2.2 m MPG telescope at ESO La Silla Obser-vatory in Chile. The images were reduced with the GRONDpipeline (Kr¨uhler et al. 2008), which applies de-bias and flat-field corrections, stacks images and provides astrometry cali-bration. Due to the bright host galaxy we disabled line by linefitting of the sky subtraction for the GROND NIR data sincethis caused over subtraction artifacts. Since the photometrybackground estimation is limited by the extended structureof the host galaxy and not the large-scale variation in thebackground of the image. We do not expect any adverse ef-fects from this change. Unfiltered imaging of AT 2016jbu wasalso obtained by B. Monard. Observations of AT 2016jbu https://ned.ipac.caltech.edu/ MNRAS , 1–28 (2021)
S. J. Brennan et al. were taken at the Kleinkaroo Observatory (KKO), Calitz-dorp (Western Cape, South Africa) using a 30cm telescopeMeade RCX400 f/8 and CCD camera SBIG ST8-XME in 2 × r -band sequence stars. Nightly images resulted from stacking(typically 5 to 8) individual images.We also recovered a number of archival images coveringthe site of AT 2016jbu. Two epochs of g and r imaging fromthe Dark Energy Camera (DECam) (Flaugher et al. 2015)mounted on the 4 m Blanco Telescope at the Cerro TololoInter-American Observatory (CTIO) were obtained from theNOIRLab Astro Data Archive. The science-ready reduced“InstCal” images were used in our analysis. In addition tothese, we downloaded deep imaging taken in 2005 with theMOSAIC-II imager (the previous camera on the 4 m BlancoTelescope). As for the DECam data, the “InstCal” reductionsof MOSAIC-II images were used. We note that the filtersused for the MOSAIC-II images (Harris V and R , Washing-ton C Harris & Canterna 1979) are different from the restof our archival dataset. The Harris filters were calibrated toJohnson-Cousins V and R . The Washington C filter datais more problematic, as this bandpass lies between Johnson-Cousins U and B . We calibrated our photometry to the latter,but this should be interpreted with appropriate caution.Deep Very Large Telescope (VLT) + OmegaCAM imagestaken with i , g , and r filters in 2013, 2014, and 2015, respec-tively, were downloaded from the ESO archive. The WideField Imager (WFI) mounted on the 2.2-m MPG telescopeat La Silla also observed NGC 2442 on a number of occasionsbetween 1999 and 2010 in B , V , and R ; these images are ofparticular interest as they are quite deep, and extend ourmonitoring of the progenitor as far back as −
15 years. Boththe OmegaCAM and WFI data were reduced using standardprocedures in
IRAF .NED contains a number of historical images of NGC 2442,dating back to 1978. We examined each of these but foundnone that contained a credible source at the position ofAT 2016jbu.Several transient surveys also provided photometric mea-surements for AT 2016jbu. Gaia G -band photometry forAT 2016jbu was downloaded from the Gaia Science Alertsweb pages. As this photometry was taken with a broad filterthat covers approximately V and R , we did not attempt tocalibrate it onto the standard system. V -band imaging wasalso taken as part of the All-Sky Automated Survey for Su-pernovae (ASAS-SN Shappee et al. 2014; Kochanek et al.2017) .The OGLE IV Transient Detection System (Koz(cid:32)lowskiet al. 2013; Wyrzykowski et al. 2014) also identifiedAT 2016jbu, and reported I -band photometry via the OGLEwebpages .The Panchromatic Robotic Optical Monitoring and Po-larimetry Telescopes (PROMPT) (Reichart et al. 2005) ob-tained imaging of AT 2016jbu in BV RI filters; and as dis- IRAF is distributed by the National Optical Astronomy Obser-vatory, which is operated by the Association of Universities for Re-search in Astronomy (AURA) under cooperative agreement withthe National Science Foundation http://ogle.astrouw.edu.pl/ogle4/transients/ cussed in Sect. 5.1.1, unfiltered PROMPT observations ofNGC 2442 were also used to constrain the activity of theprogenitor of AT 2016jbu over the preceding decade. Imageswere taken with the PROMPT1, PROMPT3, PROMPT4,PROMPT6, PROMPT7 and PROMPT8 robotic telescopes(all located at the CTIO). PROMPT4 and PROMPT6 havea diameter of 40 cm while PROMPT1, PROMPT3 andPROMPT8 have a diameter of 60 cm and PROMPT7 hasa diameter of 80 cm. All images collected with the PROMPTunits were dark subtracted and flat-field corrected. In casemultiple images were taken in consecutive exposures, theframes were registered and stacked to produce a single image.NGC 2442 was also serendipitously observed with the FO-cal Reducer/low dispersion Spectrograph 2 (FORS2) as partof the late-time follow-up campaign for SN 2015F (Cartieret al. 2017). Unfortunately, most of these data were takenwith relatively long exposures, and AT 2016jbu was satu-rated. However, a number of pre-discovery images from thesecond half of 2016, as well as late time images from 2018 areof use. These data were reduced (bias subtraction and flatfielding) using standard iraf tasks. UV and optical imaging was obtained with the
Neil GehrelsSwift Observatory ( Swift ) with the Ultra-Violet Optical Tele-scope (UVOT). The pipeline reduced data was downloadedfrom the
Swift
Data Center. The photometric reduction fol-lows the same basic outline as Brown et al. (2009). In short,a 5 (cid:48)(cid:48) radius aperture is used to measure the counts for the co-incidence loss correction, a 3 (cid:48)(cid:48) source aperture (based on theerror) was used for the aperture photometry and applying anaperture correction as appropriate (based on the average PSFin the
Swift
HEASARC’s calibration database (CALDB) andzeropoints from Breeveld et al. (2011).Subsequent to the photometric reduction of our
Swift data,there was an update to the
Swift
CALDB with time depen-dant zero-points which we have not accounted for. Given thatour
Swift observations occurred in early 2017, this wouldamount to a ∼
3% shift in zero-point and would not leadto a significant change in our lightcurve.
Near-infrared imaging was obtained with NTT+SOFI as partof the ePESSTO survey, and with GROND as mentionedpreviously. In both cases
JHK/Ks filters were used. SOFIdata were reduced using the PESSTO pipeline (Smartt et al.2015). Data were corrected for flat-field and illumination,sky subtraction was performed using (in most instances) off-target dithers, before individual frames were co-added tomake a science-ready image.In addition to the follow-up data obtained for AT 2016jbuwith SOFI, we examined pre-discovery SOFI images taken aspart of the PESSTO follow-up campaign for SN 2015F. Wedownloaded reduced images from the ESO Phase 3 archivewhich covered the period up to April 2014. Two subsequentepochs of SOFI imaging from 2016 Oct were taken afterPESSTO SSDR3 was released, and so we downloaded theraw data from the ESO archive, and reduced these using thePESSTO pipeline as for the rest of the SOFI follow-up imag-ing.
MNRAS000
MNRAS000 , 1–28 (2021) hotometric and spectroscopic evolution of AT 2016jbu Fortuitously, the ESO VISTA telescope equipped withVIRCAM observed NGC 2442 as part of the VISTA Hemi-sphere Survey (VHS) in Dec. 2016. We downloaded the re-duced images as part of the ESO Phase 3 data release fromVHS via ESO Science Portal. Photometry was performed us-ing
AutoPhOT , see Sect. 2.6.
We queried the WISE data archive at the NASA/IPAC in-frared science archive, and found that AT 2016jbu was ob-served in the course of the NEOWISE reactivation mission(Mainzer et al. 2014). As the spatial resolution of
WISE is lowcompared to our other imaging, we were careful to select onlysources that were spatially coincident with the position ofAT 2016jbu. There were numerous detections of AT 2016jbuin the W W Event B (MJD 57784 . ± . Ch
1, although there is a more point-like source present in Ch
2. No point source is seen in Ch3 and Ch4. K18 reportvalues of 0 . ± . . ± . Ch Ch Ch Ch A target of opportunity observation (ObsID: 0794580101)was obtained with XMM-Newton (Jansen et al. 2001) on 2017Jan 26 (MJD 57779) for a duration of ∼
57 ks. The data fromEPIC-PN (Str¨uder et al. 2001) were analysed using the lat-est version of the Science Analysis Software, SASv18 includ-ing the most updated calibration files. The source and back-ground were extracted from a 15 (cid:48)(cid:48) region avoiding a brightnearby source. Standard filtering and screening criteria werethen applied to create the final products.X-ray imaging was also taken with the XRT on board Swift .These observations are much less sensitive than the XMM-Newton data, and so we do not expect a detection. Usingthe online XRT analysis tools (Evans et al. 2007, 2009) weco-added all XRT images covering the site of AT 2016jbuavailable in the Swift data archive. No source was detectedcoincident with AT 2016jbu in the resulting ∼
100 ks stackedimage.
AutoPHoT pipeline
The dataset presented in this paper for AT 2016jbu com-prises approximately ∼ http://xmm.esac.esa.int/sas/ and hetrogeneous datasets, we have developed a new photo-metric pipeline called AutoPhOT ( AUTOmated PHotom-etry Of Transients ; Brennan et al. in prep).
AutoPhOT has been used to measure all photometry presented in this pa-per, with the exception of imaging from space telescopes (i.e.
Swift , Gaia , WISE , Spitzer , XMM-Newton OM and
HST ), aswell as from ground based surveys which have custom pho-tometric pipelines (i.e. ASAS-SN and OGLE).
AutoPhOT is a Python3 -based photometry pipelinebuilt on a number of commonly used astronomy packages,mostly from astropy . AutoPhOT is able to handle hetro-geneous data from different telescopes, and performs all stepsnecessary to produce a science-ready lightcurve with minimaluser interaction.In brief,
AutoPhOT will build a model for the PointSpread Function (PSF) in an image from bright isolatedsources in the field (if no suitable sources are present then
Au-toPhOT will fall back to aperture photometry). This PSFis then fitted to the transient to measure the instrumentalmagnitude. To calibrate the instrumental magnitude onto thestandard system (either AB magnitudes for Sloan-like filtersor Vega magnitudes for Johnson-Cousins filters) for this workon AT 2016jbu, the zeropoint for each image is found fromcatalogued standards in the field. For griz filters, the zero-point was calculated from magnitudes of sources in the fieldtaken from the SkyMapper Southern Survey (Onken et al.2019). For Johnson-Cousins filters, we used the tertiary stan-dards in NGC 2442 presented by Pastorello et al. (2009). Inthe case of the NIR data (
JHK ) we used sources taken fromthe Two Micron All Sky Survey (2MASS; Skrutskie et al.2006). There is no u -band photometry covering this portionof the sky. We use U -band photometry from Cartier et al.(2017) and convert to u -band using Table 1 in Jester et al.(2005). We include Swope photometry from K18 in Fig. 2 toshow that our u -band is consistent. AutoPhOT utilises a local version of
Astrometry.net (Barron et al. 2008) for astrometric calibration when imageastrometric calibration meta-data is missing or incorrect. Ininstances where AT 2016jbu could not be clearly detectedin an image, AutoPhOT performs template subtraction us-ing hotpants (Becker 2015), before doing forced photom-etry at the location of AT 2016jbu. Based on the results ofthis, we report either a magnitude or a 3 σ upper limit to themagnitude of AT 2016jbu. Artificial sources of comparablemagnitude were injected and recovered to confirm these mea-surements and to determine realistic uncertainties, account-ing for the local background and the presence of additionalcorrelated noise resulting from the template subtraction.Finally, in order to remove cases where a poor subtractionleads to spurious detections, we require that the FWHM ofany detected source agrees with the FWHM measured forthe image to within one pixel, as well as being above ourcalculated limiting magnitude. In practice we find these aregood acceptance tests to avoid false positives, especially inthe pre-discovery lightcurve of AT 2016jbu.We present the observed lightcurve of AT 2016jbu in Fig. 2, https://github.com/Astro-Sean/autophot https://anaconda.org/astro-sean/autophot http://astrometry.net/ https://github.com/acbecker/hotpants MNRAS , 1–28 (2021)
S. J. Brennan et al. and show a portion of the tables of calibrated photometryin Appendix B (the full tables are presented in the onlinesupplementary materials).
Most of our spectroscopic monitoring of AT 2016jbu was ob-tained with NTT+EFOSC2 through the ePESSTO collabo-ration. With the exception of the first classification spectrumreported by Fraser et al. (2017), observations were taken withgrisms Gr ∼
390 and R ∼ ∼ few ˚A, so that the [O i ] λ α where Gr (Valenti et al. 2014).The automatic reduction pipeline splits the first and secondorder spectra into red and blue arms and rectifies them usinga Legrendre Polynomial. Data is then trimmed, flat-fieldedusing images taken during the observing block and cleaned ofcosmic rays. Red and blue arms are then flux and wavelengthcalibrated and then merged into a 1D spectrum.A single spectrum was obtained with the WiFeS IFU spec-trograph, mounted on the ANU 2.3m telescope. This spec-trum was reduced with the PyWiFeS pipeline (Childresset al. 2014).All optical spectra are listed in Table 1 and are shown inFig. 7. For completeness, we also include the classificationspectrum of AT 2016jbu in our analysis obtained with the duPont 2.5-m telescope + WFCCD (and reported in Bose et al.2017), as it is the earliest spectrum available of the transient,see also Fig. 3.We present a single NIR spectrum taken in the low-dispersion and high-throughput prism mode with FIRE (Sim-coe et al. 2013) mounted on one of the twin Magellan Tele-scopes (Fig. 16). The spectrum was obtained using the ABBA“nod-along-the-slit” technique at the parallactic angle. Foursets of ABBA dithers totalling 16 individual frames and2028.8s of on-target integration time were obtained. Detailsof the reduction and telluric correction process are outlinedby Hsiao et al. (2019).In addition, we present two spectra taken with GeminiSouth + Flamingos2 (Eikenberry et al. 2006) in long-slitmode. An ABBA dither pattern was used for observationsof both AT 2016jbu and a telluric standard. These data were https://github.com/LCOGT/floyds_pipeline reduced using the gemini.f2 package within iraf . A prelimi-nary flux calibration was made using the telluric standard oneach night (in both cases a Vega analog was observed), andthis was then adjusted slightly to match the J − H colour ofAT 2016jbu from contemporaneous NIR imaging. Swift +UVOT spectra were reduced using the uvotpypython package (Kuin 2014) and calibrations from Kuinet al. (2015).
We present our complete lightcurve for AT 2016jbu in Fig. 2,spanning from ∼
10 years before maximum brightness (MJD:57784.4) to ∼ r -band) will be referred to as “ Event A ”, and thesubsequent brighter peak is “
Event B ”. Event A is first de-tected around three months (phase: −
91 d) before the
EventB maximum in VLT+FORS2 imaging (Fraser et al. 2017).Phases presented in this paper for AT 2016jbu and otherSN 2009ip-like transients will always be in reference to
EventB maximum light. The rise and decline of this first peak isclearly seen in r -band and sparsely sampled by Gaia in G -band. Event A has a rise time to peak of ∼
60 days, reachingan apparent magnitude r ∼ .
12 mag (absolute magni-tude − .
96 mag). We then see a short decline in r -bandfor ∼ Event B .The second event has a faster rise time of ∼
19 days, peak-ing at r ∼ − .
26 mag). After ∼
20 days past the
Event B maximum, a flattening is seen inSloan- gri and Cousins BV that persists for ∼ ∼ .
04 mag d − . At ∼
50 days, a rapid drop isseen at optical wavelengths, with the drop being more pro-nounced in the redder bands and less in the bluer bands.After the drop there is a second flattening. After two monthsfrom the
Event B peak, the optical bands flatten out witha decay of ∼ .
015 mag d − and remain this way until theseasonal gap at ∼
120 days.Late time bumps and undulations in the lightcurves of SNeare commonly associated with late time CSM interaction,when SN ejecta collide with dense stratified and/or clumpyCSM far away from the progenitor, providing a source of latetime energy (Fox et al. 2013; Martin et al. 2015; Arcavi et al.2017; Nyholm et al. 2017; Andrews & Smith 2018; Moriyaet al. 2020).A re-brightening event is seen after ∼
120 days and is seenclearly in
BVGgr -bands. We miss the initial rebrighteningevent in our ground-based data, so it is unclear if this isa plateau lasting across the seasonal gap or a rebrighteningevent. However, evidence for a rebrightening in the lightcurveis seen in
Gaia - G (See Fig. 2). We can deduce that thisevent occurred between +160 and +195 days from our Gaia - G data, where we have G = 18 .
69 mag at +160 days, but anincrease to 18.12 mag one month later. An additional bumpis seen in
Gaia - G at +345 days. We observe G = 18 .
95 magat +316 d and G = 18 .
88 mag at +342 d before AT 2016jbufades to G = 19 .
72 mag a month later.
MNRAS , 1–28 (2021) hotometric and spectroscopic evolution of AT 2016jbu A pp a r e n t m a g n i t u d e + o ff s e t R i s e D e c li n e P l a t e a u K n ee A n k l e Event A Event B W W K -18.0 H -16.5 J -15.0 z -13.5 I -12.0 i -10.5 G -9.0 R -7.5 r -6.0 V -4.5 V UVOT -3.0 g -1.5 B +0.0 B UVOT +1.5 U +3.0 U UVOT +4.5 u +6.0 UVW
UVM
UVW W W K -18.0 H -16.5 J -15.0 z -13.5 I -12.0 i -10.5 G -9.0 R -7.5 r -6.0 V -4.5 V UVOT -3.0 g -1.5 B +0.0 B UVOT +1.5 U +3.0 U UVOT +4.5 u +6.0 UVW
UVM
UVW A pp a r e n t m a g n i t u d e *No offset included in this panelVISTA+VIRCAMGRONDWFIDECam OGLESwift+UVOTVLT+FORS2HST KKOXMMASAS-SNNTT+SOFI LCOGTWatcherSwope NTT+EFOSC2PromptWISE VST+OMEGAcamGaiaCTIO+MOSAIC150 125 100 75 50 25 0 25 50 75 100 125 200 300 400 500 600 700 8006500 6000 5500 5000 4500 4000 3500 3000 2500 2000 1500 1000 500150 100 50 0 50 100 200 400 600 8006000 5000 4000 3000 2000 1000Days since V-band maximum 1510505101520253035161820222426Days since V-band maximum Figure 2.
The complete multi-band observed photometry for AT 2016jbu. The upper panel covers the period from the start of
EventA (First detection at −
91 d from VLT+FORS2) until the end of our monitoring campaign ∼ Event B peak. Offsets (listedin the legend) have been applied to each filter for clarity in the upper panels only. Note that there is a change in scale in the X-axisafter 135 days. We indicate
Event A and the rise and decline of the peak of
Event B . Epochs where spectra were taken are marked withvertical ticks. We also include the published Swope photometry from K18 (given as filled circles) to demonstrate that our photometry isconsistent. We include a horizontal magenta dotted line in all panels to demonstrate the early 2019 F W magnitudes (Paper II). Weonly plot error bars greater than 0.1 mag. The lower panel shows detections and upper limits over a period from ∼
18 years prior to
EventA . No offsets are included in this panel; light points with arrows show upper limits, while solid points are detections.MNRAS , 1–28 (2021)
S. J. Brennan et al. Å ] F Na I D Ca II NIRHHH H Fe II (42)
Figure 3.
Classification spectrum of AT 2016jbu obtained with the Du Pont 2.5-m telescope and WFCCD (and reported in Bose et al.2017) taken on 2016 December 31 ( − . Event A .The green dashed line is the blackbody fit with T BB ∼ α and H β dominate the spectra and are both well fitted with a P Cygniprofile with an additional emission component. We can also distinguish the Na I D lines superimposed on He i λ ii λλ ∼ −
700 km s − . A noise spike at5397 ˚A has been removed manually. For the purpose of discussion, we adopt the nomenclaturefor features seen in the lightcurve of SN 2009ip from Grahamet al. (2014); rise, decline, knee, and ankle. We do not desig-nate a “bump” phase as while SN 2009ip shows a clear bumpat ∼
20 d, this is not seen in AT 2016jbu. The rise begins at ∼ + 22 days prior to V -band maximum. The decline phasebegins at V -band maximum. The plateau begins at ∼ + 20days, when the decline gradient flattens out initially. The knee stage is ∼ + 45 days past maximum when a sharp drop isseen in the lightcurve, and the ankle is the flattening of thelightcurve after ∼
65 days before the seasonal gap.
There exists a growing sample of SN 2009ip-like transientswhich evolve almost identically in terms of their photometryand spectroscopy, in the years prior to, and during their mainluminous events.We focus on a small sample of objects that show com-mon similarities to AT 2016jbu. For the purpose of a qual-itative study, we will compare AT 2016jbu with SN 2009ip(Fraser et al. 2013; Graham et al. 2014), SN 2015bh (Elias-Rosa et al. 2016; Th¨one et al. 2017), LSQ13zm (Tartagliaet al. 2016), SN 2013gc (Reguitti et al. 2019) and SN 2016bdu(Pastorello et al. 2018). We will refer to these transients (in-cluding AT 2016jbu) as SN 2009ip-like transients.We also include SN 1996al (Benetti et al. 2016) in ourSN 2009ip-like sample. Although no pre-explosion variabil-ity or an
Event A/B lightcurve was detected, SN 1996alshows a similar bumpy decay from maximum and a simi-lar spectral evolution ( see Fig. 17, 19, and 20). SN 1996alshows no sign of explosively nucleosynthesized material; e.g.[O i ] λλ , Ni mass are simi-lar to what is found for AT 2016jbu and other SN 2009ip-liketransients, see Paper II. Benetti et al. (2016) suggest that thisis consistent with a fall-back supernova in a highly structured environment, and we discuss this possibility for AT 2016jbuin Paper II.We will also discuss SN 2018cnf (Pastorello et al. 2019);a previously classified Type IIn SN (Prentice et al. 2018).Although Pastorello et al. (2019) argues that SN 2018cnfdisplays many of the characteristics of SN 2009ip, it doesnot show the degree of asymmetry in H α when compared toAT 2016jbu but does show pre-explosion variability and gen-eral spectral evolution similar to SN 2009ip-like transients.Fig. 4 shows that all these transients show a relatively slowcolor evolution, typically seen in Type IIn SNe (Taddia et al.2013; Nyholm et al. 2020a). Where color information is avail-able, SN 2009ip-like transients initially appear red ∼ B − V ) for AT 2016jbu,SN 2015bh and SN 2009ip. These three transients span col-ors from ( B − V ) ∼ ∼ −
20 d to ∼ ∼ −
10 d.In general, after the peak of
Event B the transients begin tocool and again evolve towards the red.For the first ∼
20 days after
Event B , AT 2016jbu followsthe trend of other transients, which is seen clearly in ( U − B ) ,( B − V ) , ( g − r ) , and ( r − i ) . At ∼
20 d AT 2016jbu flattensin ( U − B ) and ( r − i ) , similar to SN 1996al and SN 2018cnf,whereas SN 2009ip flattens at ∼
40 d in ( U − B ) . This phasecorresponds with the plateau stage in AT 2016jbu. This fea-ture is also seen in ( r − i ) and ( u − g ) , where AT 2016jbuplateaus at ∼
20 d and then slowly evolves to the blue.This behaviour is also seen in ( B − V ) and ( g − r ) , wherea color change is observed at ∼
50 d, followed by AT 2016jburemaining at approximately constant color until the seasonalgap at ∼
120 d. SN 2018cnf follows the trend of AT 2016jbuquite closely in ( B − V ) but this abrupt transition to the blueis seen at ∼
30 d in SN 2018cnf, and ∼
60 d in AT 2016jbu.AT 2016jbu and SN 2018cnf are distinct in their ( g − r ) evo-lution, as they match SN 2009ip and SN 2016bdu closely until ∼
50 d, after which AT 2016jbu remains at an approximatelyconstant color, while SN 2009ip and SN 2016bdu make anabrupt shift to the red.
MNRAS , 1–28 (2021) hotometric and spectroscopic evolution of AT 2016jbu Table 1.
Log of optical, UV, and NIR spectra obtained for AT 2016jbu. MJD refers to the start of the exposure. Phase is with respectto the time of V -band maximum (MJD 57784 . ± . − . − . − . − . − . − . − . − . − . − . − . Swift + UVOT UV Grism2017-01-26 57779.3 − . − . − . − . − . − . ∗ FTS+FLOYDS red2017-02-14 57798.5 +14.1 FTS+FLOYDS red/blue2017-02-17 57801.5 +17.1 FTS+FLOYDS red/blue2017-02-19 57803.2 +18.8 NTT+EFOSC2 Gr ∗ FTS+FLOYDS red/blue2017-03-11 57823.5 +39.1 ∗ FTS+FLOYDS red2017-03-24 57836.0 +51.6 NTT+EFOSC2 Gr ∗ NTT+EFOSC2 Gr * Spectrum not plotted in Fig. 2 due to low S/N but still used in analysis when applicablefor Fig. 9.
We note that filters that cover the wavelength of H α (viz. r,V ) show the color change at ∼
60 d in AT 2016jbu (i.e( B − V ) , ( g − r ) , and ( r − i ) ), whereas those that do notcover H α show a similar feature at ∼
30 d i.e. ( U − B ) and( u − g ) .( B − V ) , ( g − r ) , and ( r − i ) do not show this trend butrather a transition to the blue at ∼
60 d. This correspondsto the
Knee/Ankle stage in the lightcurve. At this time wesee an increase in the relative strength of the H α blue shoul-der emission component, see Sect. 4.1. The undulations seenin these color indices may indicate part of the circumstellar environment becoming transparent through geometric dilu-tion of opaque material or sufficient cooling. This may beevident from spectral changes mentioned in Sect. 3, such asthe narrowing of H α , as well as the lack of significant velocityevolution after ∼
60 d, see Fig. 9.AT 2016jbu shows ( u − g ) ∼ − . EventB . The blue appearance of the transient suggest low values ofextinction in the immediate environment of the explosion siteand the line-of-sight. K18 discuss the presence of pre-existingdust surrounding the progenitor site and likely present during
Event A/B in December 2016. One way of explaining the blue
MNRAS , 1–28 (2021) S. J. Brennan et al. -1.00.0 ( U B ) o AT 2016jbu SN 2009ip SN 2015bh SN 2016bdu SN 1996al SN 2013gc LSQ13zm SN 2018cnf -1.00.00.01.0 ( B V ) o ( u g ) o -1.00.01.00.01.0 ( g r ) o ( r i ) o R i s e D e c li n e P l a t e a u K n ee A n k l e
200 300 400 -1.00.01.0* rescaled axisDays since Event B maximum
Figure 4.
Intrinsic color evolution of AT 2016jbu and SN 2009ip-like transients. All transients have been corrected for extinction usingthe values from Table. 2. X-axis gives days from
Event B maximum light. We include a broken X-axis to exclude the seasonal gap forAT 2016jbu. Data shown for AT 2016jbu has been regrouped into 1 day bins and weighted averaged. Error bars are shown for all objects,and we do not plot any point with an uncertainty greater than 0.5 mag. The different stages of evolution of AT 2016jbu are marked withgrey dashed vertical bands.
Table 2.
Properties of SN 2009ip-like transient events. Values reported are used consistently throughout this work. The time of peak iswith respect to the
Event B maximum. Where quoted, Ni masses are upper limits.Transient z A V [mag] µ [mag] Peak (MJD) Ni [M (cid:12) ] ReferenceAT 2016jbu 0.00489 0.556 31.60 57784 ≤ ≤ ≤ ≤ Galactic Extinction only. If A V not mentioned in reference, we take values from NED. With respect to
Event B maximum light in V -band. color of AT 2016jbu as well as the inferred presence of dust(which should make the transient appear red) could be a disk-like CSM geometry, with low density CSM along the poles,and that is partially face on. A trait of SN 2009ip-like transients is erratic photometricvariability in the period leading up to
Event A and
Event B . SN 2009ip itself underwent a series of outbursts over a threeyear period before its
Event A/B in 2012.The lower panel of Fig. 2 shows all pre-
Event A/B ob-servations for AT 2016jbu from ground based instruments.The majority of these observations are from the PROMPTtelescope array, and have been host subtracted using latetime r -band templates from EFOSC2. Unfortunately, theseimages are relatively shallow. In addition, we recovered sev-eral images from the LCO network which were obtained forthe follow-up campaign of SN 2015F (Cartier et al. 2017). MNRAS000
Event A/B ob-servations for AT 2016jbu from ground based instruments.The majority of these observations are from the PROMPTtelescope array, and have been host subtracted using latetime r -band templates from EFOSC2. Unfortunately, theseimages are relatively shallow. In addition, we recovered sev-eral images from the LCO network which were obtained forthe follow-up campaign of SN 2015F (Cartier et al. 2017). MNRAS000 , 1–28 (2021) hotometric and spectroscopic evolution of AT 2016jbu ADate: 2012-02-11Telescope: Prompt M r : 20.78S/N: 9 BDate: 2013-05-10Telescope: Prompt M r : 20.32S/N: 5CDate: 2015-02-16Telescope: Prompt M r : 20.21S/N: 9 DDate: 2012-12-16Telescope: Prompt M r : 20.14S/N: 19 Figure 5.
Sample of pre-explosion detections from PROMPTat the progenitor location. Center of cutout corresponds toAT 2016jbu progenitor location. Red circle signifies aperture withradius 1 . × FWHM placed in the center of the cutout. As men-tioned in Sect. 2.1, these unfiltered images have been host sub-tracted using r band templates. Template subtractions performedusing AutoPhOT and
HOTPANTS (Becker 2015), see Sect. 2.6.
These images have been host subtracted using templatesfrom LCO taken in 2019. We also present several imagestaken VLT+OMEGAcam which are deeper than our tem-plates and are hence not host subtracted. For completenesswe also plot detections of the progenitor of AT 2016jbu from
HST in Fig. 2, which we discuss in Paper II.If AT 2016jbu underwent a similar series of outbursts priorto
Event A/B as seen in other SN 2009ip-like transients,then we would expect to only detect the brightest of these.SN 2009ip experiences variability at least three years prior toits main events.For AT 2016jbu, several significant detections are foundwith r ∼
20 mag in the years prior to
Event A/B . For ouradopted distance modulus and extinction parameters, thesedetections correspond to an absolute magnitude of M r ∼− . R ∼ − . ∼
450 d days prior to their
Event A/B . TheAT 2016jbu progenitor is seen in
HST images around −
400 dshowing clear variations. A single DECam image in r bandgives a detection at r ∼ . ± .
26 mag at −
352 d whichroughly agrees with our
F350LP lightcurve at this time (ifwe presume H α is the dominant contributor to the flux). Wepresent and further discuss HST detections in Paper II.We note that we detect a point source at the site of AT 2016jbu in several PROMPT images but not in anyof the LCO, WFI, NTT+EFOSC2/SOFI, OmegaCAM orVISTA+VIRCAM pre-explosion images. However, a clear de-tection is made with CTIO+DECAM that is compatible withour
HST observations (see Paper II for more discussion ofthis).In Fig. 5 we show a selection of cutouts from our hostsubtracted PROMPT images, showing the region aroundAT 2016jbu. While some of the detections that
AutoPhOT recovers are marginal, others are quite clearly detected, andso we are confident that the pre-discovery variability is real. Ifthese are indeed genuine detections, then AT 2016jbu is pos-sibly undergoing rapid variability similar to SN 2009ip andSN 2015bh in the years leading up to their
Event A . The highcadence of our PROMPT imaging and the inclusion of H α inthe Lum filter plausibly explain why we have not detected theprogenitor in outburst in data from any other instrument.AT 2016jbu never displays the magnitude of pre-explosionvariability that is seen in other SN 2009ip-like transients. Ourhigh cadence observations from PROMPT could detect anysimilarly bright event in the prior years. It is unlikely thatwe missed any such eruption, especially in the last 1.5 yearsuntil maximum light. AT 2016jbu could have undergone lessextreme variability, that may be due to a different progenitoror progenitor environment.AT 2016jbu could be undergoing a slow rise up until thebeginning of
Event A similar to UGC 2773-OT (Smith et al.2016) (Intriguingly this is also seen in Luminous Red No-vae, Pastorello et al. 2021; Williams et al. 2015 - we returnto this in Paper II). Fitting a linear rise to the PROMPTpre-explosion detections (i.e. excluding the
HST and DECamdetections) gives a slope of − . ± × − mag d − and in-tercept of 19 . ± .
19 mag. If we extrapolate this line fit to −
60 d (roughly the beginning of r -band coverage for EventA ) we find a value of r extrapolate ∼ .
11 mag which is verysimilar to the detected magnitude at −
59 d of r ∼ .
09 mag.However, this is speculative, and accounting for the sporadicdetections in the preceding years, and the non-detections indeeper images e.g. from LCO see lower panel of Fig. 2, it ismore likely that AT 2016jbu is undergoing rapid variability(similar to SN 2009ip) which is serendipitously detected inour PROMPT images due to their high cadence.
Fig. 6 shows
Swift +UVOT observations around maximumlight. All bands show a sharp increase at ∼ −
18 d, consistentwith our optical lightcurve. The
Swift +UVOT can constrainthe initial
Event B rise to some time between ∼ − . ∼ − . UVW2 shows a pos-sible bump beginning at ∼
24 d that spans a few days. Thisbump is also evident in
UVM2 at the same time. This bumpis consistent with the emergence of a blue shoulder emissionin H α (See Sect. 4.1) and it is possible that we are seeing aninteraction site between ejecta and CSM at this time. No clear X-ray source was found consistent with the locationof AT 2016jbu in the XMM data taken at -5 d. Using the
MNRAS , 1–28 (2021) S. J. Brennan et al.
30 20 10 0 10 20 30 40 50Days since V-Band Maximum10121416182022 O b s e r v e d m a g n i t u d e + c o n s t a n t UVW2-4UVM2-3UVW1-2 U_UVOT+0B_UVOT+1V_UVOT+2UVW2-4UVM2-3UVW1-2 U_UVOT+0B_UVOT+1V_UVOT+2
Figure 6.
Swift + UVOT lightcurve for AT 2016jbu. All pho-tometry is host subtracted. Offsets are given in the legend anduncertainties are plotted for all points. sosta tool on the data from the PN camera we obtain a 3 σ upper limit of < . × − counts s − for AT 2016jbu; whilethe summed MOS1+MOS2 data gives a limit of < . × − counts s − . Assuming a photon index of 2, the upper limitto the observed flux in the 0.2–10 keV energy range is 1 . × − erg cm − s − .For comparison, SN 2009ip was detected in X-rays in the0.3–10 keV energy band with a flux of (1 . ± . × − ergcm − s − , as well as having an upper limit on its hard X-rayflux around optical maximum (Margutti et al. 2014).X-ray observations can tell us about the ejecta-CSM inter-action as well as the medium into which they are expand-ing into (Dwarkadas & Gruszko 2012). The non-detectionfor AT 2016jbu provides little information on the nature of Event A/B . Making a qualitative comparison to SN 2009ipwe note that AT 2016jbu is not as X-ray bright, and thismay reflect different explosion energies, CSM environmentsor line-of-sight effects.
We measure fluxes for AT 2016jbu in
Spitzer
IRAC Ch ± Ch ± Ch Ch Ch Ch HST observations,K18 finds that the progenitor of AT 2016jbu is consistent withthe progenitor system having a significant IR excess from arelatively compact, dusty shell. The dust mass in the im-mediate environment of the progenitor system is small (a few × − M (cid:12) ). However, the different epochs of the HST (takenin 2016) and
Spitzer (taken in 2003) data suggest they maybe at different phases of evolution. Fig. 2 shows that the siteof AT 2016jbu underwent multiple outbursts between 2006and 2013, and, as mentioned by K18, fitting a single SED tothe
HST and
Spitzer datasets may be somewhat misleading.
We present our high cadence spectral coverage of AT 2016jbuin Fig. 7. Our spectra begin at −
31 days and show an initialappearance similar to a Type IIn SN, i.e. narrow emissionfeatures seen in H and a blue continuum. Our first spec-tra coincide with the approximate peak of
Event A . Afteraround a week, additional absorption and emission featuresemerge in the Balmer series, which we illustrate in Fig. 8and plot the evolution of in Fig. 9. The spectrum does notvary significantly over the first month of evolution aside fromthe continuum becoming progressively bluer with time. H α shows a P Cygni profile with an emission component withFWHM ∼ − and a blue shifted absorption com-ponent with a minimum at ∼ −
600 km s − . The narrowemission lines likely arise from an unshocked CSM environ-ment around the progenitor. Over time AT 2016jbu developsa multi-component emission profile seen clearly in H α thatpersists until late times. We do not find any clear signs of ex-plosively nucleosynthesised material at late times, and indeedthe spectral evolution appears to be dominated by CSM in-teraction at all times. We discuss the evolution of the Balmerseries in Sect. 4.1. In Sect. 4.2 we discuss the evolution ofCa ii features and model late time emission profiles. Sect. 4.3discusses the evolution of several isolated, strong iron lines.Sect. 4.4 discusses the evolution of He i emission and makesqualitative comparisons between He i features and the opticallightcurve. We present UV and NIR spectra in Sect. 4.5 andSect. 4.6 respectively. The most prominent spectral features are the Balmer lines,which show dramatic evolution over time. In particular theH α profile, which shows a complex, multi-component evo-lution, provides insight to the CSM environment, mass-losshistory and explosion sequence. Although SN 2009ip neverdisplayed obvious multi-component emission features, a red-shoulder emission is seen at late times (Fraser et al. 2013).We present the evolution of H α for AT 2016jbu at severalepochs showing the major changes in Fig. 8.We conducted spectral decomposition to understand lineshape and the ejecta-CSM interaction. We used a MarkovChain Monte Carlo (MCMC) approach for fitting a multi-component spectral profile (Newville et al. 2014) using a cus-tom python3 script. When fitting, absorption componentsare constrained to be blueward of the rest wavelength of eachline to reflect a P Cygni absorption. All lines are fitted over asmall wavelength window and we include a pseudo-continuumduring our fitting, which is allowed to vary. Fitting the H α evolution is performed on each spectrum consecutively, usingthe fitted parameters from the previous model as the start-ing guess for the next. This is reset after the observing gapat +202 days. Fig. 8 presents fitted models to the H α profileat epochs where significant change are seen. The FWHM andpeak wavelength for H α are illustrated in Fig. 9. Days − to −
25: Our first spectrum coincides with theapproximate peak of
Event A (Fig. 2). H α can be mod-elled by a P Cygni profile with an absorption minimum at ∼ −
700 km s − superimposed on a broad component at ∼ + 700 km s − with a FWHM of ∼ − . This canbe interpreted as a narrow P Cygni with extended, electron- MNRAS000
700 km s − superimposed on a broad component at ∼ + 700 km s − with a FWHM of ∼ − . This canbe interpreted as a narrow P Cygni with extended, electron- MNRAS000 , 1–28 (2021) hotometric and spectroscopic evolution of AT 2016jbu Å ] L o g F [ e r g s c m Å ] + o ff s e t d d d d * d d d d d d d d d d +2 d +4 d +7 d +8 d +14 d +17 d +19 d +20 d +24 d +25 d +27 d +34 d +52 d +56 d +60 d +73 d * +90 d +121 d +202 d +203 d +241 d +270 d +299 d +346 d +384 d +419 d Na DHe I Ca NIR[O I] O I[Ca II]H H H HFe II (42) He I
Figure 7.
Spectral evolution of AT 2016jbu. Wavelength given in rest frame. Flux given in log scale. Prominent spectral lines and strongabsorption bands are labelled. Colors instruments used (see Tab. 1); black: NTT+EFOSC2, blue: FTS+FLOYDS, red: WiFeS, green:DuPont. Spectra marked with an asterisk have been smoothed using a Gaussian filter of FWHM 1 ˚A. scattering wings, as often seen in Type IIn SN spectra (seereview by Filippenko 1997).
Days − to +4: We see a gradual decay in amplitudeof the core broad emission until the H α emission is bestfit by an intermediate width Lorentzian profile (FWHM ∼ − ) and P Cygni absorption. We see ourLorentzian profile has broad wings, possibly due to electronscattering along the line-of-sight (Chugai 2001). At −
14 days,a blue broad absorption component clearly emerges at ∼ -5000 km s − with an initial FWHM of ∼ − . The trough of this absorption features slows to ∼ -3200 km s − at+3 d. Panel B in Fig. 8, shows H α at +1 days with a strongLorentzian emission with the now obvious blue absorption.This feature indicates that there is fast moving material thatwas not seen in the initial spectra. Assuming free expansion,we set an upper limit on the distance travelled by this ma-terial to ∼ . × cm. A similar feature was also seen inSN 2009ip, (e.g. Fig. 2 of Fraser et al. 2013) around the EventB maximum. A persistent second absorption feature was alsoseen in SN 2015bh (Elias-Rosa et al. 2016) which remained in
MNRAS , 1–28 (2021) S. J. Brennan et al.
Phase: 31.4 d A Phase: 0.8 d B F Phase: +18.8 d C Phase: +89.7 d D6400 6450 6500 6550 6600 6650 6700Wavelength [ Å ]Phase: +298.9 d E10 5 0 5 10Velocity [10 km s ] Figure 8.
Multi-component evolution of H α over a period of ∼ < + 120 d), and Gaussian emission andabsorption thereafter. Epochs are given in each panel, lines arecoloured such that yellow = core emission, red = redshifted emis-sion, green = P-Cygni absorption, cyan = high velocity absorptionand blue is blueshifted emission. In panel A an additional emissioncomponent could be included to account for the blue excess shown,although this can simply be extended electron scattering wings. absorption until several weeks after the Event B maximum,when it was replaced by an emission feature at approximatelythe same velocity. As shown in panel B of Fig. 8, the fastestmaterial is moving at ∼ − , as measured by the ex-treme blue edge of this absorption feature. This high velocityis suggestive of fast moving material ejected in a CCSN (Fil-ippenko 1997). The broad Gaussian component seen earlierhas diminished in intensity and a single Lorentzian emissionfits the H α well. This broad component may have followedquickly expanding material from the transient outburst ejectasometime around the beginning on Event A , (which at thisstage may now have become optically thin or may have beenovertaken by the faster material), whereas the remaining in-termediate emission profile is powered by shock interaction.
50 0 50 100 150 200 250 300 350 4000123456 A b s . V e l o c i t y [ k m s ] Red
Blue abs
Core P-Cygni
Blue em
50 0 50 100 150 200 250 300 350 400Days since V-Band Maximum0123456 F W H M [ k m s ] Figure 9.
Evolution of fitted parameters for H α . The upper panelshows the absolute velocity evolution of each feature. We fit apower decay law with index 0.4 dex to the blue emission fromwhen it first appears ( ∼ +18d) until the seasonal gap ( ∼ +125 d)indicated by the blue dashed line. The is also fitted for the redshoulder emission (with a different normalisation constant) as thered dashed line. We include a purple dotted line at 1200 km s − that matches the late time red and blue emission components. Thelower panel shows the FWHM evolution of each of the components.We do not plot the redshifted broad emission fitted during the firstthree epochs in either panel. Days +7 to +34: A persistent P Cygni profile is still seenbut a dramatic change is seen in the overall H α profile, nowbeing dominated by a red-shifted broad Gaussian feature cen-tered at ∼ +2200 km s − and FWHM ∼ − . Theblue absorption component has now vanished and been re-placed with an emission profile with a slightly lower velocity, − − at +18 d, seen in panel C of Fig. 8. Overthe following month, this line moves towards slower veloci-ties with a decreasing FWHM. The blue shoulder emission isclearly seen at ∼ +18 d and remains roughly constant in am-plitude (with respect to the core component) until ∼ +34 d.At +34 d this line now has a FWHM ∼ − . By+52 d this blue emission line has grown considerably in ampli-tude with respect to the core component. During this periodthe relative strength of the red and blue component beginsto change, indicating on-going interaction and/or changingopacities. We note that prior to +52 d, this H α profile maybe fitted with a single, broad emission component with a PCygni profile. However, during our fitting a significant blueexcess was always present during +7 d to +34 d. Allowingfor both a blue and red emission component during thesetimes, allows each consistent component to evolve smoothlyinto later spectra, as is seen in Fig 8 and Fig. 9. Days +52 to +120: H α shows a symmetric double-peakedemission profile. The blue emission feature grows in ampli-tude with respect to the core and red emission profile, andmoves red-ward from ∼ − − to ∼ − − .Over the two month span, the red emission component MNRAS000
Evolution of fitted parameters for H α . The upper panelshows the absolute velocity evolution of each feature. We fit apower decay law with index 0.4 dex to the blue emission fromwhen it first appears ( ∼ +18d) until the seasonal gap ( ∼ +125 d)indicated by the blue dashed line. The is also fitted for the redshoulder emission (with a different normalisation constant) as thered dashed line. We include a purple dotted line at 1200 km s − that matches the late time red and blue emission components. Thelower panel shows the FWHM evolution of each of the components.We do not plot the redshifted broad emission fitted during the firstthree epochs in either panel. Days +7 to +34: A persistent P Cygni profile is still seenbut a dramatic change is seen in the overall H α profile, nowbeing dominated by a red-shifted broad Gaussian feature cen-tered at ∼ +2200 km s − and FWHM ∼ − . Theblue absorption component has now vanished and been re-placed with an emission profile with a slightly lower velocity, − − at +18 d, seen in panel C of Fig. 8. Overthe following month, this line moves towards slower veloci-ties with a decreasing FWHM. The blue shoulder emission isclearly seen at ∼ +18 d and remains roughly constant in am-plitude (with respect to the core component) until ∼ +34 d.At +34 d this line now has a FWHM ∼ − . By+52 d this blue emission line has grown considerably in ampli-tude with respect to the core component. During this periodthe relative strength of the red and blue component beginsto change, indicating on-going interaction and/or changingopacities. We note that prior to +52 d, this H α profile maybe fitted with a single, broad emission component with a PCygni profile. However, during our fitting a significant blueexcess was always present during +7 d to +34 d. Allowingfor both a blue and red emission component during thesetimes, allows each consistent component to evolve smoothlyinto later spectra, as is seen in Fig 8 and Fig. 9. Days +52 to +120: H α shows a symmetric double-peakedemission profile. The blue emission feature grows in ampli-tude with respect to the core and red emission profile, andmoves red-ward from ∼ − − to ∼ − − .Over the two month span, the red emission component MNRAS000 , 1–28 (2021) hotometric and spectroscopic evolution of AT 2016jbu Å ]0.00.20.40.60.81.0 N o r m a li s e d F -31.4d Blue emission at +89.7d +89.7d12 9 6 3 0 3 6 9 12Velocity [10 km s ] Figure 10. H α profile at −
31 d (red) and +90 d (green) forAT 2016jbu. The +90 d profile has had a strong blue emissionprofile (given by dotted blue line) subtracted and we plot the resid-ual in green. Each spectra is normalised at 6563 ˚A. The profile at+90 d has been blue-shifted by 4˚A ( ∼ −
180 km s − ) to match thepeak at the H α rest wavelength (6563 ˚A) of the profile at −
31 d. changes from FWHM of ∼ − to ∼ − and its peak velocity shifts by ∼
200 km s − , from ∼ +1000 km s − to ∼ − . We note that the profileat +90 d is similar to our earliest spectra, with the excep-tion of a prominent blue component, see Fig. 10. The earliestprofile of H α at −
31 d is reminiscent of some stages duringan eruptive outburst from a massive star (for example Var C;Humphreys et al. 2014). We plot the profile of the +90 d pro-file in Fig. 10 with a blue-shifted Lorentzian profile removed.The profiles are very similar in overall shape with a slightlybroader red-core component in the +90 d spectrum. A possi-ble interpretation is the P Cygni-like profile seen in our −
31 dspectra is associated with the events during/causing
Event A (for example a stellar merger or eruptive outburst) and theblue side emission is associated with events during/causing
Event B (for example a core-collapse or CSM interaction).
Days +203 to +420: The red and blue components nowhave similar FWHM of ∼ − and ∼ − respectively. The overall H α profile is an almost symmetricdouble peak as shown in panel D of Fig. 8. The height ofthese peaks vary slightly with respect to each-other duringthis timeframe. After this time we no longer fit a P Cygniabsorption profile, and our spectra can be fitted well usingthree emission components. We justify this as any opaquematerial may have become optically thin after ∼ α is seen for the remainder of our observations. The threeemission profiles remain at their respective wavelengths andthe approximate same width. The overall evolution of H α suggests that AT 2016jbu underwent a large mass loss event(whether that be a SN or extreme mass loss episode) in ahighly aspherical environment. This ejected mass then inter-acts with dense CSM forming a multi-component H α profileas well as a bumpy lightcurve. F Å ] F
15 10 5 0 5 10 15Velocity [10 km s ] Figure 11.
Calcium NIR triplet fit for +345 d. The individualcomponents of the primary Ca ii NIR triplet is given by the bluedashed lines in both plots. The upper panel shows the emissionprofile with the inclusion of O i λ ii NIR triplet emission shown in purple (Region B). BothO i λ Sect. 4.1 indicates that AT 2016jbu has a highly non-sphericalenvironment. We investigate similar trends in other emis-sion profiles. We explore the Ca ii NIR triplet λλλ ii NIR triplet appears in emission at approximately the sametime as blue-shifted emission in H α ( ∼ +18 d) and at earlytimes shows P Cygni absorption minima at velocities similarto H α . For profile fitting, the wavelength separation betweenthe three components of the NIR triplet was held fixed, whilethe three components were also constrained to have the sameFHWM. Amplitude ratios between the three lines were con-strained to physically plausible values between the opticallythin and optically thick regimes (Herbig & Soderblom 1980).The Ca ii NIR triplet is seen in our −
31 d and −
26 d spectra,although at low S/N, and appears to show P Cygni absorp-tions with minima at ∼ −
650 km s − (however model fittingreturns large, inconsistent errors).The Ca ii NIR triplet (re-)emerges around +20 d and isfitted well with three Lorentzian components with FWHM of4000 km s − and velocity of +700 km s − with a P Cygnilike absorption with a minimum at ∼ +500 km s − , similarto that seen in H α until day +120 d. This would suggestthis emission is from a single (symmetric) emitting locationwith some opaque material present. After ∼
200 days theCa ii NIR triplet appears to narrow. At +204 d, fitting themultiplet alone gives FWHM ∼ − and a velocity MNRAS , 1–28 (2021) S. J. Brennan et al. of +500 km s − . Assuming a single emitting region howeverdoes not provide a satisfactory fit, resulting in a noticeableexcess towards the blue as can be seen in Fig. 11.We explore two scenarios for the Ca ii NIR triplet evo-lution after +200 d. In the first, we assume that the Ca ii emission comes from the same regions as H α (as suggestedin Sect. 4.1) i.e two spatially separated emitting regions. Weallow the first region to be fitted with the above restrictions(fixed line separation, single common FWHM), we refer tothis as Region A. A second kinematically distinct multipletis added (we refer to this as Region B) and simultaneouslyfitted with additional constraints that the lines have the sameFWHM as the region A and the amplitude ratio of the Ca ii NIR triplet being emitted from region B is some multiple ofthe region A. Region B represents this blue-shifted materialalso seen in H α . The second scenario has an additional Gaus-sian representing O i λ ii emitting region.As shown in Fig. 11 both scenarios give an acceptablefit to spectrum at +345 d. Fitting a single Gaussian emis-sion line representing O i λ ≈ − redshifted by ∼
800 km s − . Al-ternatively adding an additional Ca ii emission profile wefind a good fit at FWHM ≈ − and blue-shifted by ∼ − − . Although the fitting is inconclusive, thisdoes not exclude a complex asymmetrical CSM structure pro-ducing these multiple emitting regions along the line-of-sight.Although both scenarios give reasonable fits, the FWHMand velocities deduced for both scenarios are not seen else-where in the spectrum at +345 d. It is possible that theregion(s) producing the Ca ii NIR triplet is separated from Hemitting areas although detailed modelling is needed to con-firm. We note however one should expect a similar flux fromO i λ i λ ii NIR triplet and [Ca ii ] having a broadened appearance com-pared to earlier spectra. This may indicate an increase in thevelocity of the region where these lines form, similar to whatis seen in H α in Sect. 4.1. As temperatures and opacities drop the spectra of many CC-SNe become dominated by iron lines, as well as Na i and Ca ii .We notice persistent permitted Fe group transitions through-out the evolution of AT 2016jbu which is likely pre-existingiron in the progenitor envelope. Our initial spectra display theFe ii λλλ , , −
31 d we measure the absorption minimum ofFe ii multiplet 42 at −
750 km s − . This is the same velocityas the fitted absorption profile from H α /H β see Fig. 8A. Wecan assume that this lines originate in similar regions.The Fe ii multiplet 42 appears in our late time spectra, seeFig. 12. Fe ii lines in general appear with P Cygni profiles atlate times. It is difficult to measure the absorption minimumof the Fe ii profile due to severe blending. However, usingseveral relatively isolated Fe ii lines at +345 d we measurean absorption minimum of ∼ − − . The values is similar to the velocity offset for the red and blue emissioncomponents seen in H α . This suggest that these lines areoriginating in the same region. We see ionised He i , detected sporadically over the ∼ . i lines dis-play the degree of asymmetry seen in hydrogen, althoughtransients exist displaying double-peaked helium lines, suchas the Type Ibn SN 2006jc; (Foley et al. 2007; Pastorello et al.2008), as well as some displaying asymmetric He i and sym-metric H emission e.g. the Type Ibn/IIb SN 2018gjx (Prenticeet al. 2020).We show the evolution of He i λ i λ −
31 d spec-tra. He i λ −
14 d with a boxy pro-file that is poorly fit with a single Lorentzian emission line.He i λ ∼ − . Note the blue absorption featurein H α is also first seen at this time. The line begins to broadenover the next month, peaking at FWHM ∼ − at ∼ + 28 d. After +51 d, He i λ i λ ∼ − centered at restwavelength. We see this same FWHM in the red and blueshoulders in H α (Sect. 4.1). We find that a single emissionprofile matches the He i λ α we alsofind that He i λ ∼ ±
400 km s − from their rest wavelength, andeach has a FWHM of ∼ − . Unlike H α , no thirdcore emission component is needed. Low resolution spectrapreclude further investigation, but if He i λ α . An increase in the strength of He i was alsoseen in the Type IIn SN 1996al and was interpreted as asignature of strengthening CSM interaction (Benetti et al.2016).We plot the evolution of He i λ λλ −
31 d spec-trum there is a clear P Cygni profile centered at 5898 ˚A. Theemission is likely caused by Na I D with the possibility ofsome absorption contamination from He i λ ∼ −
450 km s − with respect to 5890 ˚A.This is similar to H α . At −
13 d, He i λ i λ i λ −
14 d with anabsorption trough at ∼ −
500 km s − , similar to H α . Af-ter the seasonal gap He i λ ∼ − . This is likely dominatedby Na I D with minor contamination from He i λ MNRAS , 1–28 (2021) hotometric and spectroscopic evolution of AT 2016jbu Å ] L o g F O I O I O I O I H e I H e I H e I H e I H e I H e I HHHHH N I Na D M g I ] [ F e II ][ S II ] [ O II ][ O I ] [ C a II ] C a II F e I Fe IIFe IIFe II F e II F e II F e II F e II Fe II Paschen O I Ca II NIR H e I He I [N II]
Figure 12.
Optical spectrum of AT 2016jbu at +345 d from NTT+EFOSC2. We include labels for strong emission features with redvertical lines. Several forbidden transitions are marked in blue. Double-peaked H α is clearly visible and dominates the spectra. He i λ i λ Å ] L o g F [ e r g s c m Å ] + o ff s e t +418d+383d+345d+298d+269d+240d+201d+120d+89d+80d+51d+33d+26d+24d+18d+6d+4d-3d-4d-5d-11d-13d-14d-25d-27d-31d9000 6000 3000 0 3000 6000 9000 12000Velocity [ km s ] 6900 7000 7100 7200 7300Wavelength [ Å ] L o g F [ e r g s c m Å ] + o ff s e t +418d+383d+345d+298d+269d+240d+202d+201d+120d+89d+80d+51d+33d+26d+24d+18d+6d+4d-3d-4d-5d-11d-13d-14d-25d-27d-31d6000 3000 0 3000 6000 9000Velocity [ km s ] Figure 13.
Evolution of He i λ i λ i lines (5876 ˚A and 7065 ˚A) are marked with a vertical line, while Na I D λλ , i lines is given in the upper axis. Each spectrum has been normalised to a peak value of unity. We plot the evolution of the pseudo-Equivalent Width(pEW) (a pseudo-continuum is fitted over a small wavelengthwindow) of the two seemingly isolated He i λλ , i λλ , ∼
120 days. After the seasonal gap, bothemission lines increase dramatically in pEW, until ∼ +300 dafter which the pEW declines. A similar jump in He i was seenin SN 1996al (Benetti et al. 2016). This decline coincides withthe narrowing and increase in amplitude of the blue, red, andcore emission components of H α . He i emission is expected to be formed in the de-excitation/recombination region of the shock wave (Cheva-lier & Kirshner 1978; Gillet & Fokin 2014). As mentioned inSect. 4.1, after ∼ α grows in amplitude and narrows considerably, likely due tochanging opacities. This jump in pEW may represent a timewhen shocked material is no longer obscured and photons canescape freely from the interaction sites. We reach a similarconclusion for He i . If the trend in both He i lines is linked tothe H α emitting regions, then it is likely that the late timeHe i might also be double-peaked.Fig. 2 shows a rebrightening/flattening after the seasonal MNRAS , 1–28 (2021) S. J. Brennan et al.
50 0 50 100 150 200 250 300 350 400Days since Event B maximum10 p E W [ Å ] R i s e D e c li n e P l a t e a u K n ee A n k l e He I
He I H ( Blue ) H ( Red ) H ( Core ) Figure 14.
Evolution of pEW for He i λ , λ α components. The He i emission appears to be roughly constantuntil the knee / ankle stage when it increases rapidly. After ∼ +300 d the pEW of He i again begins to decrease. The measurementof pEW is based on a single emission component fit which providesa reasonable fit at late times. He i λ t <
220 ddue to its low pEW and contamination from H α . gap. This is seen best in Gaia - G . The trend seen in He i λ λ We present a single UV spectum in Fig. 15 taken with
Swift +UVOT on 2017 January 22. The spectrum has quitelow S/N towards the red with a very tenuous detection of theBalmer series. It is likely that λ > λ < ∼ − ) emission line is thestrongest feature seen. It centered at ∼ α aroundthis time. We are unsure of the identification of this emissionline, however there is a strong Fe ii line at ∼ ii line here and no other emission features at com-parable strength. Swift observations of SN 2009ip do showthis emission line but it is much weaker than that seen inAT 2016jbu (Margutti et al. 2014). This particular emissionline has been seen in several Type IIP SNe with UV coveragesuch as SN 1999em and SN 2005cs (see Gal-Yam et al. 2008,and references therein). However, the Type IIP SNe discussedby Gal-Yam et al. (2008) also show strong emission fromMg ii λ ∼ − − which is likely due to Mg ii λ We present our NIR spectra in Fig. 16 covering the peak of
Event A as well as the rise and peak of
Event B . Pa β λ α , with a strong blue absorp-tion profile that is not present in the −
31 d FIRE spectrumbut which appears in the FLAMINGOS-2 −
12 d spectra. Atthis phase the blue absorption is already seen in H α and H β .Pa β is also broader at −
31 d and narrows at −
12 d, similarto the H α evolution shown in Fig. 8 at −
31 d and +1 d.There is a strong He i λ γ . At −
31 d this line appears in absorption at rest wavelength,while by −
12 d the line is in emission. This helium featuremay be thermally excited and this is supported by the black-body temperature seen peaking at this time (see Paper II).We see an absorption trough bluewards of λ γ λ β ). There appears to be a flux excess beyond 2.1 µ min the FIRE spectrum at −
31 d. This may represent emissionfrom a CO bandhead, possibly signifying some pre-existingdust during
Event A . However, the S/N is extremely low inthis region of the spectrum (see the grey shaded region inFig. 16), and it is likely that the apparent “excess” is due tobright K -band sky contamination rather than CO emission. We will discuss AT 2016jbu and their relation to SN 2009ip-like objects, mainly their photometric similarities inSect. 5.1.1 and their spectroscopic appearance in Sect. 5.1.2,in particular the appearance of their H α emission pro-files is varies times during their evolution (Sect. 5.1.3). Wewill briefly address the most controversial topic of whetherAT 2016jbu, and subsequently all SN 2009ip-like objects, aregenuine CCSNe in Sect. 5.2, although we will return to thisin greater detail in Paper II. For this paper we focus the discussion on the photometric andspectral comparison between AT 2016jbu and similar tran-sients. In Paper II we discuss topics including the progenitorof AT 2016jbu using pre-explosion images, the environmentaround the progenitor and a non-terminal explosion scenario.We adopt the basic parameters in Table. 2 in our analysis ofSN 2009ip-like transients.
We compare the R / r -band lightcurves of a sample ofSN 2009ip-like transients events in Fig. 17. In cases where r -band photometry was not available, Johnson-Cousin R -bandis shown. The adopted extinction and distance moduli aregiven in Tab. 2. The photometric evolution for SN 2009ip-like transients is undoubtedly similar. Our sample of tran-sients all show a series of outbursts in the years prior to Event A , as seen in Fig. 17. AT 2016jbu shows several cleardetections within ∼
10 years before the peak of
Event B . Sim-ilar outbursts are seen in other SN 2009ip-like transients (seeFig. 17).The duration of
Event A varies between each transient.
MNRAS , 1–28 (2021) hotometric and spectroscopic evolution of AT 2016jbu Å ] L o g F [ e r g s c m Å ] Fe II Fe II Fe II Fe II Fe II Fe II Fe II Fe II Mg II Ca II Figure 15.
Swift + UVOT spectrum for AT 2016jbu taken on 2017 January 22 (MJD: 57775, Phase: −
18 d). Wavelength in given inrest frame and the spectrum is corrected for Galactic extinction ( A V = 0.556 mag). The spectrum is given in black with the grey shadedregion showing the uncertainty. Paschen Brackettt1 1.2 1.4 1.6 1.8 2 2.2Wavelength [ m ] L o g F [ e r g s c m Å ] He I He I
FLAMINGOS Phase: 10.7 d FLAMINGOS Phase: +1.9 d FIRE Phase: 28.0 d Figure 16.
NIR spectra of AT 2016jbu, covering the peak of
Event A as well as the rise and peak of
Event B . H and He i are clearlyseen in all spectra. The FIRE spectrum (blue) has been smoothed for presentation and shows what appears to be an excess redwards of2 . µm . This excess is likely due to spectra being saturated by the bright K -band sky. A b s o l u t e M a g n i t u d e AT 2016jbuSN 2009ip SN 2015bhSN 2016bdu SN 1996alSN 2013gc LSQ13zmSN 2018cnfAT 2016jbuSN 2009ip SN 2015bhSN 2016bdu SN 1996alSN 2013gc LSQ13zmSN 2018cnf
Figure 17.
Pre-explosion outbursts and the main luminous event for the sample of SN 2009ip-like transients. SN 2009ip (Sloan r) is takenfrom Fraser et al. (2013); Graham et al. (2014), SN 2015bh ( R ) from Th¨one et al. (2017), SN 2016bdu ( r ) and SN 2013gc (R) taken fromReguitti et al. (2019), SN 1996al ( R ) from Benetti et al. (2016), SN 2018cnf ( r ) from Pastorello et al. (2019) and LSQ13zm ( R ) is takenfrom Tartaglia et al. (2016). All data given in Vega magnitudes (Blanton & Roweis 2007). We do not show limiting magnitudes in thisfigure for clarity. All events show an initial rise to a magnitude of ∼ -14 (if coverage available) followed by a second rise to ∼ −
18 roughly30 days later. Our sample of SN 2009ip-like transients all show outbursts in the months/years prior to their luminous events.MNRAS , 1–28 (2021) S. J. Brennan et al.
100 0 100 200 300 400Days since Event B maximum1817161514131211 A b s o l u t e M a g n i t u d e AT 2016jbuSN 2009ipSN 2015bhSN 2016bdu SN 1996alSN 2013gc LSQ13zmSN 2018cnfAT 2016jbuSN 2009ipSN 2015bhSN 2016bdu SN 1996alSN 2013gc LSQ13zmSN 2018cnf
Figure 18.
Same as Fig. 17, but focusing
Event A/B . AllSN 2009ip-like transients show a similar
Event B (lightcurve),although
Event A tends to be more diverse (if observations areavailable). AT 2016jbu shows a major rebrightening after ∼
200 ddays not seen in other SN 2009ip-like transients.
For SN 2009ip,
Event A lasts for ∼ ∼ − . ∼ − . ∼
17 days tomaximum in
Event B to ∼ − ± . − . Event B peak is still within a standarddeviation of this.Curiously, several of the transients in our sample show theirfirst initial bump around the same time, approximately 20days post maximum; see Fig. 18. AT 2016jbu shows no majorbumps in its lightcurve, but instead flattens slightly, whereasSN 2009ip and SN 2018cnf show a clear and prominent bumpat ∼
20 d.From ∼
60 – 120 d, AT 2016jbu appears to follow the ex-trapolated decline of SN 2009ip (see Fig. 18). However, whenAT 2016jbu emerged from behind the Sun at +200 d, itshowed a large increase in magnitude in all bands. No otherSN 2009ip-like transient showed a comparable behaviour.At ∼
200 d, AT 2016jbu is almost 1 mag brighter thanSN 2009ip. As we see a change in He i pEW around this time(Sect. 4.4), we suggest that this change reflects an increasein CSM interaction. This is not clearly seen in Hα at thistime and may reflect a complex CSM environment. We notethat SN 2013gc may also show similar signs of interaction at ∼
200 d.
The spectra of SN 2009ip-like transients remain remarkablysimilar as they evolve. Fig. 19 shows our sample of extinctioncorrected SN 2009ip-like transients at several phases duringtheir evolution. All objects initially appear similar to a clas-sical Type IIn SN, with T BB ∼ , K and prominentnarrow lines seen in the Balmer series, an indication that allobjects are interacting with dense slow-moving CSM. Severalweeks after maximum brightness prominent emission linesare seen, mainly in He i , Ca ii and H. Several Fe ii lines are clearly seen. It is at this time that asymmetries begin toevolve in H α which persist until very late times. We dis-cuss the evolution of H α in Sect. 5.1.3. After ∼ α continues to dominate the optical spectra. H β isstrong in SN 2009ip, SN 2015bh and SN 2016bdu but notice-able weaker in AT 2016jbu. At late times there are strikinglysimilar emission lines seen in all SN 2009ip-like transients α comparison We show a zoom in on H α in Fig. 20, where the spectra areplotted in order of “ double-peaked ”-ness i.e according to thelevel of double-peaked nature of the H α line profile. We arbi-trarily define double-peaked -ness as the strength and separa-tion between the two emission peaks (if any) seen in H α . Allobjects also appear to show an additional high velocity blueabsorption in their Balmer lines, as seen in panel B of Fig. 20.Note that the spectroscopic data for SN 1996al only beginsat 22 days past Event B , when we can already see the emer-gence of a broad blue component. At intermediate times, ∼ ∼
10 months, all transients now show multi-component profiles. Each transient displays different velocityand FWHM values for their red and blue components. ForSN 2009ip, Fraser et al. (2015) notes a red component at+500 km s − at late times, this shoulder is also seen in H β .We measure the same component at +625 km s − while fit-ting for an additional blue component at −
510 km s − . Ourfit is illustrated in Fig. 21. In the case of SN 2009ip, thisred shoulder only appeared at ∼ − − and +900 km s − respectively and areseen as early as ∼ α profile.We include a close up of the H α profile of η Car in Fig. 20,based on VLT+MUSE observations taken on 2014 Nov. 13.This spectrum was extracted from spaxels with a 14 (cid:48)(cid:48) radiusof η Car after masking nearby stars. η Car displays a multi-peaked H α profile similar to what we see in our SN 2009ip-like transients events, albeit at a lower velocity. A similarlyshaped profile is also seen in spectra obtained from lightechoes of the Great Eruption (GE) (Smith et al. 2018). Thisresemblance raises the tantalising possibility that η Car andSN 2009ip-like transients share similar progenitors or progen-itor systems.To date, it is still uncertain what caused the GE in η Car,although commonly discussed scenarios include a major erup-tion triggered by a merging event in a triple stellar system
MNRAS , 1–28 (2021) hotometric and spectroscopic evolution of AT 2016jbu +6 d +0 d d d HHHHH He I / NaD He IHe I
SN 1996alSN 2015bhSN 2016bdu AT 2016jbuSN 2009ip L o g F [ e r g s c m Å ] + o ff s e t +92 d +140 d +79 d +90 d +94 d O I O I O I O I O I O I O I H e I H e I H e I H e I H e I H e I H e I N I F e II Ca IICa II4000 4500 5000 5500 6000 6500 7000 7500 8000 8500 9000Wavelength [ Å ] +317 d +194 d +232 d +346 d +335 d Fe II (42)
Figure 19.
Spectral comparison of SN 2009ip-like transients around
Event B peak (top), three months after
Event B (middle) and latetime spectra around one year later (bottom). We include several strong Fe ii emission lines in the bottom panel as orange vertical lines.We note the remarkable similarities between AT 2016jbu and other SN 2009ip-like transients at late times. (Smith et al. 2018), mass transfer from a secondary star dur-ing periastron passages (Kashi & Soker 2010) or even a pul-sational pair-instability explosion (Woosley et al. 2007). De-spite the asymmetric H α emission lines, curiously no otherlines show such asymmetry, in particular He i . However, wecannot exclude that this is simply due to lower S/N in theseother lines, or that their lower velocities mean that any signsof asymmetry are masked by our moderate instrumental res-olution. There has been much debate on the nature of SN 2009ip-like transients (Smith & Mauerhan 2012; Fraser et al. 2013;Pastorello et al. 2013; Margutti et al. 2014; Graham et al.2014; Smith et al. 2014; Pastorello et al. 2018, 2019). Themain topic of controversy is the fate of the progenitor system.Some authors suggest that these are indeed peculiar Type IInSNe (Mauerhan et al. 2013; Prieto et al. 2013) and the activ-ity seen before any possible core-collapse (i.e the variabilityseen years prior) is due to eruptive outbursts from the pro-genitor, while
Event A and
Event B are due to interaction
MNRAS , 1–28 (2021) S. J. Brennan et al. N o r m a li s e d F + o ff s e t d d +0 d +6 d +22 d d Å ] (Rest Frame)+94 d +66 d +79 d +140 d +100 d +90 d SN 2009ip LSQ13zm SN 2016bdu SN 2015bh SN 1996al AT 2016jbu6400 6500 6600 6700+335 d +245 d +232 d +256 d +317 d +346 d Car6 3 0 3 6 6 3 0 3 6Velocity [10 km s ] 6 3 0 3 6 Figure 20. H α spectral comparison between SN 2009ip-like transients. Spectra are plotted after normalising with respect to the peak ofH α , with arbitrary flux offsets for clarity. Spectra were de-reddened using the parameters given in Table. 2. Early time spectra show aType IIn-like profile with narrow emission. while spectra ∼ α forms a double-peaked emission profile, aside from in the case of SN 2009ip (although here there is still evidence for ared shoulder component). The difference in line shape is most likely due to inclination, an idea we elaborate on in Sect. 5. We also showthe spectrum of η Car (at ∼ +150 yr) Å ] F CoreRed EmissionBlue EmissionDataModel3 2 1 0 1 2 3Velocity [10 km s ] Figure 21.
Spectral decomposition of the H α profile for SN 2009ipat +335 d. Spectra from the DEep Imaging Multi-Object Spec-trograph (DEIMOS; Faber et al. 2003), was fitted as mentionedin Fig. 8. A three-component model reproduces the observed H α profile at late times. between SN ejecta and more recently ejected CSM. The causeof the outburst immediately prior to the SN in this scenariois debated, but possible explanations include pulsational-pairinstabilities (Woosley et al. 2007), binary interaction (Kashiet al. 2013; Pastorello et al. 2019), merging of massive stars(Soker & Kashi 2013) and rapid rotation (e.g the ΩΓ-limit;Maeder & Meynet 2000). Perhaps controversially, some au-thors suggest that SN 2009ip-like events are non-terminalphenomena akin to the GE in η Car and that a survivingstar will be seen after CSM interaction fades. Whatever their true nature, Sect. 5.1 shows that AT 2016jbu displays strongsimilarities to SN 2009ip-like events, and it is likely that allthese transients share a common fate.During the 2012
Event B of SN 2009ip, material is seenmoving at ∼ − (Mauerhan et al. 2013), sug-gesting that the 2012 event was a SN. Tartaglia et al.(2016) find even faster moving material at ∼ − for LSQ13zm. We find similarly fast-moving material inAT 2016jbu ( ∼ ,
000 km s − ). However, Pastorello et al.(2013) also found evidence for ejecta moving at similar speedsin 2011 for SN 2009ip (albeit with much weaker absorptionin lines suggesting a lower mass of material). Fast materialalone cannot be used to conclusively demonstrate that theseare CCSNe.A clear sign of a terminal explosion is forbiddenemission lines from material formed during explosivenucleosynthesis/late-time stellar evolution. All CCSNe willeventually cool down significantly for the photosphere to re-cede to the innermost layers of the explosion. At this stage weexpect to see the signs of material synthesised in the explo-sion as well as material produced in the late-stages of stellarevolution such as [O i ] λλ , i ] λ Ni to Co and of Co to Fe (assuming thatCSM interaction has stopped). This stage has a characteristicdecline rate, governed by the half-life of Co ( ≈
77 days),while the luminosity is controlled by the mass of Ni ejected.We return to the mass of Ni in AT 2016jbu in Paper II. Wewill now focus on spectral signatures of synthesised material.Fig. 12 shows the late time spectra of AT 2016jbu highlight-ing prominent emission lines. Tenuous detections are made of
MNRAS , 1–28 (2021) hotometric and spectroscopic evolution of AT 2016jbu [O i ] and Mg i ], although these lines are much weaker than aretypically seen during the nebular phase of CCSNe.It is seen from late time spectra that this is still on-goingCSM interaction for AT 2016jbu, as is clear for the double-peaked H α emission. The spectra are still relatively blue (i.e.Fig. 12, λ (cid:46) In this paper, we have presented the results of our follow-up campaign for AT 2016jbu consisting of photometry up to ∼ −
31 to +420 days covering the UV, optical and NIR.We also present historical observations over the precedingdecade. In summary, the salient points of this are: • Early spectra of AT 2016jbu appear similar to a Type IInSN; showing narrow emission features and a blue con-tinuum. As the transient evolves, asymmetric featuresemerge and persist until the end of the observing cam-paign. • AT 2016jbu displays variability in the years prior to max-imum light, with outbursts reaching M r ∼ − . M r ∼ − . M r ∼ − .
26 mag, with both peaks separatedby ∼ • AT 2016jbu shows a smooth lightcurve with a majorre-brightening event occurring after the seasonal gap( ∼
200 days). An increase in He i emission is seen duringthis time, which may be a sign of increased interaction. • AT 2016jbu appears spectroscopically and photomet-rically alike to SN 2009ip, SN 2015bh, SN 2016bdu,SN 1996al, SN 2013gc and SN 2018cnf. However, the in-crease in brightness at ∼ +200 d is unique to AT 2016jbuwith respect to our sample of SN 2009ip-like transients.The color evolution is similar amongst all SN 2009ip-liketransients, while color changes can be linked with the ap-pearance of the red and blue emission components seenin H α • We compare the H α profiles of each transient and show acontinuum of asymmetry at intermediate and late times.We deduce that this may be an inclination effect and show that these emission lines can be well fitted with athree-component model. • AT 2016jbu and other SN 2009ip-like transients do notexhibit signs of explosive nucleosynthesis at late timessuch as [O i ] λλ , i ] λ DATA AVAILABILITY
The spectroscopic data underlying this article are availablein the Weizmann Interactive Supernova Data Repository at https://wiserep.weizmann.ac.il/ . The photometric dataunderlying this article are available in the article and in itsonline supplementary material.
ACKNOWLEDGEMENTS
S. J. Brennan acknowledges support from Science Founda-tion Ireland and the Royal Society (RS-EA/3471). M.F issupported by a Royal Society - Science Foundation IrelandUniversity Research Fellowship. T.M.B was funded by theCONICYT PFCHA / DOCTORADOBECAS CHILE/2017-72180113. T.W.C acknowledges the EU Funding under MarieSk(cid:32)lodowska-Curie grant H2020-MSCA-IF-2018-842471, andthanks to Thomas Kr¨uhler for GROND data reduction.M.N is supported by a Royal Astronomical Society Re-search Fellowship. B.J.S is supported by NSF grants AST-1908952, AST-1920392, AST-1911074, and NASA award80NSSC19K1717. M.S is supported by generous grants fromVillum FONDEN (13261,28021) and by a project grant(8021-00170B) from the Independent Research Fund Den-mark. L.H acknowledges support for Watcher from ScienceFoundation Ireland grant 07/RFP/PHYF295. Time domainresearch by D.J.S. is supported by NSF grants AST-1821987,1813466, & 1908972, and by the Heising-Simons Founda-tion under grant
MNRAS , 1–28 (2021) S. J. Brennan et al.
This research made use of Astropy , a community-developed core Python package for Astronomy (Astropy Col-laboration et al. 2013; Price-Whelan et al. 2018). This re-search made use of data provided by Astrometry.net Thisresearch has made use of the NASA/IPAC ExtragalacticDatabase (NED), which is operated by the Jet PropulsionLaboratory, California Institute of Technology, under con-tract with the National Aeronautics and Space Administra-tion. We acknowledge the use of public data from the Swiftdata archive. This work made use of data supplied by the UKSwift Science Data Centre at the University of Leicester. Weacknowledge Telescope Access Program (TAP) funded by theNAOC, CAS, and the Special Fund for Astronomy from theMinistry of Finance. Parts of this research were supportedby the Australian Research Council Centre of Excellence forAll Sky Astrophysics in 3 Dimensions (ASTRO 3D), throughproject number CE170100013. This work is based in part onobservations made with the Spitzer Space Telescope, whichwas operated by the Jet Propulsion Laboratory, CaliforniaInstitute of Technology under a contract with NASA. Sup-port for this work was provided by NASA through an awardissued by JPL/Caltech. This work made use of v2.2.1 of theBinary Population and Spectral Synthesis (BPASS) modelsas described in Eldridge et al. (2017) and Stanway & Eldridge(2018). This publication makes use of data products from theTwo Micron All Sky Survey, which is a joint project of theUniversity of Massachusetts and the Infrared Processing andAnalysis Center/California Institute of Technology, fundedby the National Aeronautics and Space Administration andthe National Science Foundation. This paper includes datagathered with the 6.5 meter Magellan Telescopes located atLas Campanas Observatory, Chile. This publication makesuse of data products from the Wide-field Infrared Survey Ex-plorer, which is a joint project of the University of California,Los Angeles, and the Jet Propulsion Laboratory/CaliforniaInstitute of Technology, funded by the National Aeronauticsand Space Administration. This research is based on obser-vations made with the NASA/ESA Hubble Space Telescopeobtained from the Space Telescope Science Institute, whichis operated by the Association of Universities for Research inAstronomy, Inc., under NASA contract NAS 5-26555. Theseobservations are associated with program 15645. Observa-tions were also obtained from the Hubble Legacy Archive,which is a collaboration between the Space Telescope Sci-ence Institute (STScI/NASA), the Space Telescope Euro-pean Coordinating Facility (ST-ECF/ESAC/ESA) and theCanadian Astronomy Data Centre (CADC/NRC/CSA). Thiswork made use of data from the Las Cumbres Observatorynetwork. We thank the Las Cumbres Observatory and its stafffor its continuing support of the ASAS-SN project. ASAS-SN is supported by the Gordon and Betty Moore Founda-tion through grant GBMF5490 to the Ohio State University,and NSF grants AST-1515927 and AST-1908570. Develop-ment of ASAS-SN has been supported by NSF grant AST-0908816, the Mt. Cuba Astronomical Foundation, the Centerfor Cosmology and AstroParticle Physics at the Ohio StateUniversity, the Chinese Academy of Sciences South Amer-ica Center for Astronomy (CAS- SACA), and the Villum https://astrometry.net/use.html Foundation. This research has made use of the SVO Fil-ter Profile Service supported from the Spanish MINECOthrough grant AYA2017-84089. We acknowledge ESA Gaia,DPAC and the Photometric Science Alerts Team . Thisproject used data obtained with the Dark Energy Cam-era (DECam), which was constructed by the Dark EnergySurvey (DES) collaborating institutions: Argonne NationalLab, University of California Santa Cruz, University of Cam-bridge, Centro de Investigaciones Energeticas, Medioambi-entales y Tecnologicas-Madrid, University of Chicago, Uni-versity College London, DES-Brazil consortium, Universityof Edinburgh, ETH-Zurich, University of Illinois at Urbana-Champaign, Institut de Ciencies de l’Espai, Institut de Fisicad’Altes Energies, Lawrence Berkeley National Lab, Ludwig-Maximilians Universitat, University of Michigan, NationalOptical Astronomy Observatory, University of Nottingham,Ohio State University, University of Pennsylvania, Univer-sity of Portsmouth, SLAC National Lab, Stanford University,University of Sussex, and Texas A&M University. Funding forDES, including DECam, has been provided by the U.S. De-partment of Energy, National Science Foundation, Ministryof Education and Science (Spain), Science and TechnologyFacilities Council (UK), Higher Education Funding Coun-cil (England), National Center for Supercomputing Applica-tions, Kavli Institute for Cosmological Physics, Financiadorade Estudos e Projetos, Funda¸c˜ao Carlos Chagas Filho deAmparo a Pesquisa, Conselho Nacional de DesenvolvimentoCient´ıfico e Tecnol´ogico and the Minist´erio da Ciˆencia e Tec-nologia (Brazil), the German Research Foundation-sponsoredcluster of excellence “Origin and Structure of the Universe”and the DES collaborating institutions. Based on observa-tions obtained with XMM-Newton, an ESA science mis-sion with instruments and contributions directly funded byESA Member States and NASA.Based on observations ob-tained at the international Gemini Observatory (GS-2016B-Q-22), a program of NSF’s NOIRLab, which is managedby the Association of Universities for Research in Astron-omy (AURA) under a cooperative agreement with the Na-tional Science Foundation. on behalf of the Gemini Obser-vatory partnership: the National Science Foundation (UnitedStates), National Research Council (Canada), Agencia Na-cional de Investigaci´on y Desarrollo (Chile), Ministerio deCiencia, Tecnolog´ıa e Innovaci´on (Argentina), Minist´erio daCiˆencia, Tecnologia, Inova¸c˜oes e Comunica¸c˜oes (Brazil), andKorea Astronomy and Space Science Institute (Republic ofKorea). This paper includes data gathered with the NordicOptical Telescope (PI Stritzinger) at the Observatorio delRoque de los Muchachos, La Palma, Spain. We thank Nor-bert Schartel for rapid approval and scheduling of ToO ob-servations (ObsID 0794580101). Based on observations col-lected at the European Organisation for Astronomical Re-search in the Southern Hemisphere under ESO programmes1103.D-0328, 1103.D-0328, 199.D-0143, 197.D-1075, 191.D-0935 (PESSTO, ePESSTO and ePESSTO+). In addition,data taken under ESO programmes 098.A-9099, 0104.A-9099,0100.D-0865, 098.D-0692, 100.D-0341, and 0103.D-0139(A)are presented here. Based in part on data acquired at theANU 2.3m Telescope at Siding Spring Observatory. We ac- http://svo2.cab.inta-csic.es/theory/fps/ http://gsaweb.ast.cam.ac.uk/alerts MNRAS , 1–28 (2021) hotometric and spectroscopic evolution of AT 2016jbu knowledge the traditional owners of the land on which thetelescope stands, the Gamilaraay people, and pay our re-spects to elders past, present, and emerging. Part of thefunding for GROND (both hardware as well as personnel)was generously granted from the Leibniz-Prize to Prof. G.Hasinger (DFG grant HA 1850/28-1). The OGLE project hasreceived funding from the National Science Centre, Poland,grant MAESTRO 2014/14/A/ST9/00121 to AU. /00311 andDaina No. 2017/27/L/ST9/03221. LCO data have been ob-tained via OPTICON proposals and was obtained as part ofthe Global Supernova Project. The OPTICON project hasreceived funding from the European Union’s Horizon 2020research and innovation programme under grant agreementNo 730890. REFERENCES
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APPENDIX A: AUTHOR AFFILIATIONS School of Physics, O’Brien Centre for Science North,University College Dublin, Belfield, Dublin 4, Ireland The Oskar Klein Centre, Department of Physics, AlbaNova,Stockholm University, SE-106 91 Stockholm, Sweden INAF-Osservatorio Astronomico di Padova, Vicolodell’Osservatorio 5, I-35122 Padova, Italy Department of Physics and Astronomy, University ofTurku, FI-20014, Turku, Finland The Department of Physics, The University of Auckland,Private Bag 92019, Auckland, New Zealand The Oskar Klein Centre, Department of Astronomy,AlbaNova, Stockholm University, SE-106 91 Stockholm, Sweden Max-Planck-Institut f¨ur Extraterrestrische Physik,Giessenbachstraße 1, 85748 Garching, Germany Department of Astronomy, The Ohio State University, 140W. 18th Avenue, Columbus, OH 43210, USA Center for Cosmology and AstroParticle Physics (CCAPP),The Ohio State University, 191 W. Woodruff Avenue, Colum-bus, OH 43210, USA Department of Physics and Astronomy, Texas A&MUniversity, 4242 TAMU, College Station, TX 77843, USA Cerro Tololo Inter-American Observatory, NSF’s NationalOptical-Infrared Astronomy Research Laboratory, Casilla603, La Serena, Chile Institut d’Astrophysique de Paris (IAP), CNRS & Sor-bonne Universite, France Kavli Institute for Astronomy and Astrophysics, PekingUniversity, Yi He Yuan Road 5, Hai Dian District, Beijing100871, China NINAF-Osservatorio Astronomico di Padova, Vicolodell’Osservatorio 5, I-35122 Padova, Italy Institute of Space Sciences (ICE, CSIC), Campus UAB,Carrer de Can Magrans s/n, 08193 Barcelona, Spain Center for Astrophysics | Harvard & Smithsonian, 60Garden Street, Cambridge, MA 02138, USA Department of Physics, Florida State University, 77Chieftan Way, Tallahassee, FL 32306, USA Tuorla Observatory, Department of Physics and Astron-omy, FI-20014 University of Turku, Finland. Finnish Centre for Astronomy with ESO (FINCA),FI-20014 University of Turku, Finland CBA Kleinkaroo, Calitzdorp, South Africa Departamento de Ciencias F´ısicas, Universidad AndresBello, Avda. Republica 252, Santiago, 8320000, Chile Millennium Institute of Astrophysics, Santiago, Chile Department of Astronomy/Steward Observatory, 933North Cherry Avenue, Rm. N204, Tucson, AZ 85721-0065,USA Institute for Astronomy, University of Hawai’i, 2680Woodlawn Drive, Honolulu, HI 96822, USA Astrophysics Research Centre, School of Maths andPhysics, Queen’s University Belfast, Belfast BT7 1NN, UK Mt Stromlo Observatory, The Research School of Astron-omy and Astrophysics, Australian National University, ACT2601, Australia National Centre for the Public Awareness of Science, Aus-tralian National University, Canberra, ACT 2611, Australia The ARC Centre of Excellence for All-Sky Astrophysicsin 3 Dimension (ASTRO 3D), Australia Astronomical Observatory, University of Warsaw, Al.Ujazdowskie 4, 00-478 Warszawa, Poland Unidad Mixta Internacional Franco-Chilena de As-tronom´ıa, CNRS/INSU UMI 3386 and Instituto de As-trof´ısica, Pontificia Universidad Cat´olica de Chile, Santiago,Chile Aix Marseille Univ, CNRS, CNES, LAM, Marseille,France Institute of Astronomy, Madingley Road, Cambridge,CB30HA, UK RHEA Group for ESA, European Space AstronomyCentre (ESAC-ESA), Madrid, Spain Departamento de F´ısica Te´orica y del Cosmos, Universi-dad de Granada, E-18071 Granada, Spain
MNRAS000
MNRAS000 , 1–28 (2021) hotometric and spectroscopic evolution of AT 2016jbu Tuorla Observatory, Department of Physics and Astron-omy, FI-20014 University of Turku, Finland Kavli Institute for Cosmology, Institute of Astronomy,Madingley Road, Cambridge, CB3 0HA, UK Las Cumbres Observatory, 6740 Cortona Drive, Suite 102,Goleta, CA 93117-5575, USA Department of Physics, University of California, SantaBarbara, CA 93106-9530, USA The Observatories of the Carnegie Institution for Science,813 Santa Barbara St., Pasadena, CA 91101, USA School of Physics & Astronomy, Cardiff University,Queens Buildings, The Parade, Cardiff, CF24 3AA, UK School of Physics and Astronomy, University of Southamp-ton, Southampton, Hampshire, SO17 1BJ, UK School of Physics, Trinity College Dublin, The Universityof Dublin, Dublin 2, Ireland Department of Physics, University of the Free State, POBox 339, Bloemfontein 9300, South Africa Carnegie Observatories, Las Campanas Observatory,Colina El Pino, Casilla 601, Chile Birmingham Institute for Gravitational Wave Astronomyand School of Physics and Astronomy, University of Birm-ingham, Birmingham B15 2TT, UK Institute for Astronomy, University of Edinburgh, RoyalObservatory, Blackford Hill, EH9 3HJ, UK Nucleo de Astronomıa de la Facultad de Ingenierıa yCiencias, Universidad Diego Portales, Av. Ej ´ ercito 441,Santiago, Chile Department of Physics and Astronomy University ofNorth Carolina at Chapel Hill Chapel Hill, NC 27599, USA Department of Physics, Florida State University, 77Chieftain Way, Tallahassee, FL 32306-4350, USA Department of Physics and Astronomy, Aarhus University,Ny Munkegade, DK-8000 Aarhus C, Denmark Department of Physics, University of Warwick, CoventryCV4 7 AL, UK Astrophysics Research Centre, School of Mathematics andPhysics, Queen‘s University Belfast, Belfast BT7 1NN, UK
MNRAS , 1–28 (2021) S. J. Brennan et al.
APPENDIX B: PHOTOMETRY TABLES
Table B1: Sample of full photometry table for AT 2016jbu. All mea-surements were carried out using
AutoPhoT . Phase is with respectto V -band maximum of Event B . Limiting magnitudes listed whereAT 2016jbu could not be detected, and 1 σ errors are given in parenthe-ses. UBVRIJHK filters are in Vegamags, ugriz are in AB magnitudes.Full photometry table available online.Date MJD Phase (d) u g r i z U B V R I J H K Instrument1999-12-26 51538.5 − > − > > > − > − > > − > − > > − > − > − > − > This paper has been typeset from a TEX/L A TEX file prepared by the author.MNRAS000