DRAFT VERSION JUNE 4, 2013 Preprint typeset using LATEX style emulateapj v. 5/2/11

arXiv:1306.0038v1 [astro-ph.HE] 31 May 2013

A PANCHROMATIC VIEW OF THE RESTLESS SN 2009IP REVEALS THE EXPLOSIVE EJECTION OF A MASSIVE STAR ENVELOPE
R. MARGUTTI1 , D. MILISAVLJEVIC1 , A. M. SODERBERG1 , R. CHORNOCK1 , B. A. ZAUDERER1 , K. MURASE2 , C. GUIDORZI3 , N. E. SANDERS1 , P. KUIN4 , C. FRANSSON5 , E. M. LEVESQUE6 , P. CHANDRA7 , E. BERGER1 , F. B. BIANCO8 , P. J. BROWN9 , P. CHALLIS7 , E. CHATZOPOULOS10 , C. C. CHEUNG11 , C. CHOI12 , L. CHOMIUK13,14 , N. CHUGAI15 , C. CONTRERAS16 , M. R. DROUT1 , R. FESEN17 , R. J.
FOLEY1 , W. FONG1 , A. S. FRIEDMAN1,18 , C. GALL19,20 , N. GEHRELS20 , J. HJORTH19 , E. HSIAO21 , R. KIRSHNER1 , M. IM12 , G. LELOUDAS22,19 , R. LUNNAN1 , G. H. MARION1 , J. MARTIN23 , N. MORRELL24 , K. F. NEUGENT25 , N. OMODEI26 , M. M. PHILLIPS24 , A.
REST27 , J. M. SILVERMAN10 , J. STRADER13 , M. D. STRITZINGER28 , T. SZALAI29 , N. B. UTTERBACK17 , J. VINKO29,10 , J. C. WHEELER10 , D. ARNETT30 , S. CAMPANA31 , R. CHEVALIER32 , A. GINSBURG6 , A. KAMBLE1 , P. W. A. ROMING33,34 , T. PRITCHARD34 ,
G. STRINGFELLOW6
Draft version June 4, 2013
ABSTRACT
The 2012 explosion of SN 2009ip raises questions about our understanding of the late stages of massive star evolution. Here we present a comprehensive study of SN 2009ip during its remarkable re-brightening(s). Highcadence photometric and spectroscopic observations from the GeV to the radio band obtained from a variety of ground-based and space facilities (including the VLA, Swift, Fermi, HST and XMM) constrain SN 2009ip to be a low energy (E ∼ 1050 erg for an ejecta mass ∼ 0.5 M ) and likely asymmetric explosion in a complex medium shaped by multiple eruptions of the restless progenitor star. Most of the energy is radiated as a result of the shock breaking out through a dense shell of material located at ∼ 5 × 1014 cm with M ∼ 0.1 M , ejected by the precursor outburst ∼ 40 days before the major explosion. We interpret the NIR excess of emission as signature of dust vaporization of material located further out (R > 4 × 1015 cm), the origin of which has to be connected with documented mass loss episodes in the previous years. Our modeling predicts bright neutrino emission associated with the shock break-out if the cosmic ray energy is comparable to the radiated energy. We connect this phenomenology with the explosive ejection of the outer layers of the massive progenitor star, that later interacted with material deposited in the surroundings by previous eruptions. Future observations will reveal if the luminous blue variable (LBV) progenitor star survived. Irrespective of whether the explosion was terminal, SN 2009ip brought to light the existence of new channels for sustained episodic mass-loss, the physical origin of which has yet to be identified.
Subject headings: supernovae: specific (SN 2009ip)

1 Harvard-Smithsonian Center for Astrophysics, 60 Garden St., Cambridge, MA 02138, USA.
2 Institute for Advanced Study, Princeton, New Jersey 08540, USA. 3 Department of Physics, University of Ferrara, via Saragat 1, I-44122 Ferrara, Italy. 4 University College London, MSSL, Holmbury St. Mary, Dorking, Surrey RH5 6NT, UK. 5 Department of Astronomy and the Oskar Klein Centre, Stockholm University, AlbaNova, SE-106 91 Stockholm, Sweden. 6 CASA, Department of Astrophysical and Planetary Sciences, University of Colorado, 389-UCB, Boulder, CO 80309, USA. 7 National Centre for Radio Astrophysics, Tata Institute of Fundamental Research, Pune University Campus, Ganeshkhind, Pune 411007, India. 8 Center for Cosmology and Particle Physics, New York University, 4 Washington Place, New York, NY 10003. 9 George P. and Cynthia Woods Mitchell Institute for Fundamental Physics & Astronomy, Texas A. & M. University, Department of Physics and Astronomy, 4242 TAMU, College Station, TX 77843, USA. 10 Department of Astronomy, University of Texas at Austin, Austin, TX 78712-1205, USA. 11 Space Science Division, Naval Research Laboratory, Washington, DC 20375-5352, USA. 12 CEOU/Department of Physics and Astronomy, Seoul National University, Seoul 151-742, Republic of Korea. 13 Department of Physics and Astronomy, Michigan State University, East Lansing, MI 48824, USA. 14 National Radio Astronomy Observatory, P.O. Box O, Socorro, NM 87801. 15 Institute of Astronomy, Russian Academy of Sciences, Pyatnitskaya 48, 119017, Moscow, Russian Federation. 16 Centre for Astrophysics & Supercomputing, Swinburne University of Technology, PO Box 218, Hawthorn, VIC 3122, Australia.

17 Department of Physics & Astronomy, Dartmouth College, 6127 Wilder Lab, Hanover, NH 03755, USA.
18 Massachusetts Institute of Technology, 77 Massachusetts Ave., Bldg. E51-173, Cambridge, MA 02138, USA.
19 Dark Cosmology Centre, Niels Bohr Institute, University of Copenhagen, Juliane Maries Vej 30, DK-2100 Copenhagen, Denmark.
20 NASA, Goddard Space Flight Center, 8800 Greenbelt Road, Greenbelt, MD 20771, USA.
21 Carnegie Observatories, Las Campanas Observatory, Colina El Pino, Casilla 601, Chile.
22 The Oskar Klein Centre, Department of Physics, Stockholm University, SE-10691, Stockholm, Sweden.
23 Astronomy/Physics MS HSB 314, One University Plaza Springfield, IL 62730, USA.
24 Carnegie Observatories, Las Campanas Observatory, Casilla 601, La Serena, Chile.
25 Lowell Observatory, 1400 W Mars Hill Road, Flagstaff, AZ 86001, USA.
26 W. W. Hansen Experimental Physics Laboratory, Kavli Institute for Particle Astrophysics and Cosmology, Department of Physics and SLAC National Accelerator Laboratory, Stanford University, Stanford, CA 94305, USA.
27 Space Telescope Science Institute, 3700 San Martin Dr., Baltimore, MD 21218, USA.
28 Department of Physics and Astronomy, Aarhus University, Ny Munkegade, DK-8000 Aarhus C, Denmark.
29 Department of Optics and Quantum Electronics, University of Szeged, Dóm tér 9., Szeged H-6720, Hungary.
30 Department of Astronomy and Steward Observatory, University of Arizona, Tucson, AZ 85721, USA.
31 INAF/Brera Astronomical Observatory, via Bianchi 46, 23807, Merate (LC), Italy.

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FIG. 1.— Hubble Space Telescope (HST) pre-explosion image acquired in 1999 (Program 6359; PI: Stiavelli). The location of the progenitor of SN 2009ip is marked by an arrow. SN 2009ip exploded in the outskirts of its host galaxy NGC 7259 at an angular distance of ∼ 43.4 from the host center, corresponding to ∼ 5 kpc.
1. INTRODUCTION
Standard stellar evolutionary models predict massive stars with M 40 M to spend half a Myr in the Wolf-Rayet (WR) phase before exploding as supernovae (SNe, e.g. Georgy et al. 2012 and references therein). As a result, massive stars are not expected to be H rich at the time of explosion. Yet, recent observations have questioned this picture, revealing the limitations of our current understanding of the last stages of massive star evolution and in particular the uncertainties in the commonly assumed mass loss prescriptions (e.g. Smith & Owocki 2006). Here, we present observations from an extensive, broad-band monitoring campaign of SN 2009ip (Fig. 1) during its double explosion in 2012 that revealed extreme mass-loss properties, raising questions about our understanding of the late stages of massive star evolution.
An increasingly complex picture is emerging connecting SN progenitor stars with explosion properties. The most direct link arguably comes from the detection of progenitor stars in pre-explosion images. These efforts have been successful connecting Type IIP SNe with the death of red supergiants (M ∼ 8 − 15 M , Smartt 2009). However, massive progenitor stars have proven to be more elusive (e.g. Kochanek et al. 2008): SN 2005gl constitutes the first direct evidence for a massive (M > 50 M ) and H rich star to explode as a corecollapse SN, contrary to theoretical expectations (Gal-Yam et al. 2007; Gal-Yam et al. 2009).
SN 2005gl belongs to the class of Type IIn SNe (Schlegel 1990). Their spectra show evidence for strong interaction between the explosion ejecta and a dense circumstellar medium (CSM) previously enriched by mass loss from the progenitor star. In order for the SN to appear as a Type IIn explosion, the mass loss and the core collapse have to be timed, with mass
32 Department of Astronomy, University of Virginia, P.O. Box 400325, Charlottesville, VA 22904-4325, USA.
33 Southwest Research Institute, Department of Space Science, 6220 Culebra Road, San Antonio, TX 78238, USA.
34 Department of Astronomy & Astrophysics, Penn State University, 525 Davey Lab, University Park, PA 16802, USA.

loss occurring in the decades to years before the collapse. This timing requirement constitutes a further challenge to current evolutionary models and emphasizes the importance of the progenitor mass loss in the years before the explosion in determining its observable properties.
Mass loss in massive stars can either occur through steady winds (on a typical time scale of 103 yr) or episodic outbursts lasting months to years, reminiscent of Luminous Blue Variable (LBV) eruptions (see Humphreys & Davidson 1994 for a review). SN 2005gl, with its LBV-like progenitor, established the first direct observational connection between SNe IIn and LBVs. On the other hand, there are controversial objects like SN 1961V, highlighting the present difficulty in distinguishing between a giant LBV eruption and a genuine core-collapse explosion even 50 years after the event (Van Dyk & Matheson 2012, Kochanek et al. 2011 and references therein). The dividing line between SNe and impostors can be ambiguous.
Here we report on our extensive multi-wavelength campaign to monitor the evolution of SN 2009ip, which offers an unparalleled opportunity to study the effects and causes of significant mass loss in massive stars in real time. Discovered in 2009 (Maza et al. 2009) in NGC 7259 (spiral galaxy with brightness MB ∼ −18 mag, Lauberts & Valentijn 1989), it was first mistaken as a faint SN candidate (hence the name SN 2009ip). Later observations (Miller et al. 2009; Li et al. 2009; Berger et al. 2009) showed the behavior of SN 2009ip to be consistent instead with that of LBVs. Pre-explosion Hubble Space Telescope (HST) images constrain the progenitor to be a massive star with M 60 M (Smith et al. 2010b, Foley et al. 2011), consistent with an LBV nature. The studies by Smith et al. (2010b) and Foley et al. (2011) showed that SN 2009ip underwent multiple explosions in rapid succession in 2009. Indeed, a number of LBV-like eruptions were also observed in 2010 and 2011: a detailed summary can be found in Levesque et al. (2012) and a historic light-curve is presented by Pastorello et al. (2012). Among the most important findings is the presence of blue-shifted absorption lines corresponding to ejecta traveling at a velocity of 2000−7000 km s−1 during the 2009 outbursts (Smith et al. 2010b, Foley et al. 2011), extending to v ∼ 13000 km s−1 in Semptember 2011 (Pastorello et al. 2012). Velocities this large have never been associated with LBV outbursts to date.
SN 2009ip re-brightened again on 2012 July 24 (Drake et al. 2012), only to dim considerably ∼ 40 days afterwards (hereafter referred to as the 2012a outburst). The appearance of high-velocity spectral features was first noted by Smith & Mauerhan (2012) on 2012 September 22. This was shortly followed by the major 2012 re-brightening on September 23 (2012b explosion hereafter, Brimacombe 2012, Margutti et al. 2012). SN 2009ip reached MV < −18 mag at this time, consequently questioning the actual survival of the progenitor star: SN or impostor? (Pastorello et al. 2012; Prieto et al. 2013; Mauerhan et al. 2013; Fraser et al. 2013; Soker & Kashi 2013).
Here we present a comprehensive study of SN 2009ip during its remarkable evolution in 2012. Using observations spanning more than 15 decades in wavelength, from the GeV to the radio band, we constrain the properties of the explosion and its complex environment, identify characteristic time scales that regulate the mass loss history of the progenitor star, and study the process of dust vaporization in the progenitor surroundings. We further predict the neutrino emission associated with this transient.

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This paper is organized as follows. In Section 2-6 we describe our follow up campaign and derive the observables that can be directly constrained by our data. In Section 7 we present the properties of the explosion and environment that can be inferred from the data under reasonable assumptions. In Section 8 we address the major questions raised by this explosion and speculate about answers. Conclusions are drawn in Section 9.
Uncertainties are 1σ unless stated otherwise. Following Foley et al. (2011) we adopt a distance modulus of µ = 32.05 mag corresponding to a distance dL = 24 Mpc and a Milky Way extinction E(B − V ) = 0.019 mag (Schlegel et al. 1998) with no additional host galaxy or circumstellar extinction. From VLTXshooter high-resolution spectroscopy our best estimate for the redshift of the explosion is z = 0.005720, which we adopt through out the paper. We use U, B and V for the Johnson filters. u, b and v refer to Swift-UVOT filters. Standard cosmological quantities have been adopted: H0 = 71 km s−1 Mpc−1, ΩΛ = 0.73, ΩM = 0.27. All dates are in UT and are reported with respect to MJD 56203 (2012 October 3) which corresponds to the UV peak (tpk).
2. OBSERVATIONS AND DATA ANALYSIS
Our campaign includes data from the radio band to the GeV range. We first describe the data acquisition and reduction in the UV, optical and NIR bands (which are dominated by thermal emission processes) and then describe the radio, Xray and GeV observations, which sample the portion of the spectrum where non-thermal processes are likely to dominate.
2.1. UV photometry
We initiated our Swift-UVOT (Roming et al. 2005) photometric campaign on 2009 September 10 and followed the evolution of SN 2009ip in the 6 UVOT filters up until April 2013. Swift-UVOT observations span the wavelength range
λc = 1928 Å (w2 filter) - λc = 5468 Å (v filter, central wavelength listed, see Poole et al. 2008 for details). Data have been analyzed following the prescriptions by Brown et al. (2009). We used different apertures during the fading of the 2012a outburst to maximize the signal-to-noise ratio and limit the contamination by a nearby star. For the 2012a event we used a 3 aperture for the b and v filters; 4 for the u filter and 5 for the UV filters. We correct for PSF losses following standard prescriptions. At peak SN 2009ip reaches u∼ 12.5 mag potentially at risk for major coincidence losses. For this reason we requested a smaller readout region around maximum light. Our final photometry is reported in Table 5 and shown in Fig. 2. The photometry is based on the UVOT photometric system (Poole et al. 2008) and the revised zero-points of Breeveld et al. (2011).
2.2. UV spectroscopy: Swift-UVOT and HST
Motivated by the bright UV emission and very blue colors of SN 2009ip we initiated extensive UV spectral monitoring on 2012 September 27, ∼ 6 days before maximum UV light (2012b explosion). Our campaign includes a total of 22 UVOT UV-grism low-resolution spectra and two epochs of Hubble Space Telescope (HST) observations (PI R. Kirshner), covering the period −6 days < t − tpk < +34 days.
Starting on 2012 September 27 (tpk − 6 days), a series of spectra were taken with the Swift UVOT UV grism, with a cadence of one to two days, until 2012 October 28 (tpk + 25 days), with a final long observation on 2012 November 2 (tpk + 30 days).

Details are given in Table 2. Over the course of the month the available roll angles changed. The roll angle controls the position of the grism spectrum relative to the strong zeroth orders from background stars in the grism image. The best roll angle had the spectrum lying close to the first order of some other sources, while some zeroth orders contaminated part of the spectrum in some observations. Finally, the second order overlap limits the usefulness of the red part of the UV grism spectrum. To obtain the best possible uncontaminated spectra range, the spectra were observed at a position on the detector where the second order lies next to the first order, increasing the good, uncontaminated part of the first order from about
1900 Å to 4500 Å. The spectra were extracted from the image using the
UVOTPY package. The wavelength anchor and flux calibration were the recent updates valid for locations other than the default position in the center (Kuin et al. in prep., details can be found on-line35). The spectra were extracted for a slit with the default 2.5 σ aperture and a 1 σ aperture. An aperture correction was made to the 1 σ aperture spectra which were used. The 1 σ aperture does not suffer as much from contamination as the larger aperture. Contamination from other sources and orders is readily seen when comparing the extractions of the
two apertures. The wavelength accuracy is 20 Å (1 σ), the flux accuracy (systematic) is within 10%, while the resolution is R ∼ 75 − 110 depending on the wavelength range. The error in the flux was computed from the Poisson noise in the data, as well as from the consistency of between the spectra extracted from the images on one day. The sequence of SwiftUVOT spectra is shown in Fig. 3. Figure 4 shows the UV portion of the spectra, re-normalized using the black-body fits derived in Sec. 3.
Starting from November 2012 SN 2009ip is too faint for Swift-UVOT spectroscopic observations. We continued our UV spectroscopic campaign with HST (Fig. 5). Observations with the Space Telescope Imaging Spectrograph (STIS) were taken on 2012 October 29 (tpk + 26 days) using aperture 52x0.2E1 with gratings G230LB, G430L and G750L with exposures times of 1200 s, 400 s and 100 s, respectively. The STIS 2-D images were cleaned of cosmic rays and dead pixel signatures before extraction. The extracted spectra were then matched in flux to the STSDAS/STIS pipeline 1-D data product. The spectrum is shown in Fig. 5. Further HST-COS data were acquired on 2012 November 6 (tpk + 34 days, Fig. 5, lower panel). Observations with the Cosmic Origin Spectrograph (COS) were acquired using MIRROR A + bright object aperture for 250 s. The COS data were then reprocessed with the COS calibration pipeline, CALCOS v2.13.6, and combined with the custom IDL co-addition procedure described by Danforth et al. (2010) and Shull et al. (2010). The coaddition routine interpolates all detector segments and grating settings onto a common wavelength grid, and makes a correction for the detector quantum efficiency enhancement grid. No correction for the detector hex pattern is performed. Data were obtained in four central wavelength settings in each farUV grating mode (1291, 1300, 1309, and 1318 with G130M and 1577, 1589, 1600, and 1611 with G160M) at the default focal-plane split position. The total exposure time for the farUV observation was 3100 s and 3700 s for near-UV.
2.3. Optical photometry
35 http://www.mssl.ucl.ac.uk/ npmk/Grism

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FIG. 2.— Photometric evolution of SN 2009ip in the UV, optical and NIR (filled circles). We add NIR observations of the 2012a outburst published by Pastorello et al. (2012) for t < −10 days (triangles) together with R and I-band photometry from Prieto et al. (2013) obtained during the rise-time (triangles). The shaded gray vertical bands mark the time of observed bumps in the light-curve. Our late time UVOT photometry from April 2013 is not shown here.

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FIG. 3.— Sequence of Swift-UVOT spectra of SN 2009ip covering the rise time (red to orange), peak time (shades of purple) and decay time (shades of blue) of the 2012b explosion.
Observations in the v, b and u filters were obtained with Swift-UVOT and reduced as explained in Sec. 2.1. The results from our observations are listed in Table 5. In Fig. 2 we apply a dynamical color term correction to the UVOT v, b and u filters to plot the equivalent Johnson magnitudes as obtained following the prescriptions by Poole et al. (2008). This is a minor correction to the measured magnitudes and it is not responsible for the observed light-curve bumps.
We complement our data set with R and I band photometry obtained with the UIS Barber Observatory 20-inch telescope (Pleasant Plains, IL), a 0.40 m f/6.8 refracting telescope operated by Josch Hambasch at the Remote Observatory Atacama Desert, a Celestron C9.25 operated by TG Tan (Perth, Australia), and a C11 Schmidt-Cassegrain telescope operated by Ivan Curtis (Adelaide, Australia). Exposure times ranged from 120 s to 600 s. Images were reduced following standard procedure. Each individual image in the series was measured and then averaged together over the course of the night. The brightness was measured using circular apertures adjusted for seeing conditions and sky background from an annulus set around each aperture. Twenty comparison stars within 10 of the target were selected from the AAVSO36 Photometric All-
36 http://www.aavso.org/apass.

1.4 Fe III B III/Fe III/Sc I 1.2 1.0 0.8 0.6
2000 1.4 Fe III B III/Fe III/Sc I 1.2 1.0 0.8 0.6
2000 1.4 Fe III B III/Fe III/Sc I 1.2 1.0 0.8 0.6
2000 1.4 Fe III B III/Fe III/Sc I 1.2 1.0 0.8 0.6
2000 1.4 Fe III B III/Fe III/Sc I 1.2 1.0 0.8 0.6
2000 1.4 Fe III B III/Fe III/Sc I 1.2 1.0 0.8 0.6
2000 1.4 Fe III B III/Fe III/Sc I 1.2 1.0 0.8 0.6
2000

Swift UVOT normalised spectra
Fe II Mg II Fe I

2500 Fe II

3000 Mg II Fe I

2500 Fe II

3000 Mg II Fe I

2500 Fe II

3000 Mg II Fe I

2500 Fe II

3000 Mg II Fe I

2500 Fe II

3000 Mg II Fe I

2500 Fe II

3000 Mg II Fe I

2500 3000
λ(A◦ )

∆∆∆tttpppkkk===---654ddd

2012-09-27 2012-09-28 2012-09-29

3500 4000

∆∆∆tttpppkkk===---321ddd

2012-09-30 2012-10-01 2012-10-02

3500

4000

∆ =tpk ∆ =tpk ∆ =tpk

0d 1d 2d

2012-10-03 2012-10-04 2012-10-05

3500

4000

∆ =tpk ∆ =tpk ∆ =tpk

3d 4d 5d

2012-10-06 2012-10-07 2012-10-08

3500 4000 ∆∆∆tttpppkkk===1791ddd222000111222--1-11000--1-11024

3500 4000

∆∆∆tttpppkkk===111379ddd

2012-10-16 2012-10-20 2012-10-22

3500 4000

∆∆∆tttpppkkk===223131ddd

2012-10-24 2012-10-26 2012-11-03

3500 4000

FIG. 4.— UV portion of the Swift-UVOT spectra re-normalized using the black-body fits of Sect. 3, with identifications. As time proceeds Fe III absorption features become weaker while Fe II develops stronger absorption features, consistent with the progressive decrease of the black-body temperature with time (Fig. 11).
Sky Survey. Statistical errors in the photometry for individual images were typically 0.05 magnitudes or less. Photometry taken by different telescopes on the same night are comparable within the errors. Finally, the photometry was corrected to the photometric system of Pastorello et al. (2012) using the corrections of dR = +0.046 and dI = +0.023 (Pastorello, personal communication). R and I band photometry is reported in Table 6.
A single, late-time (tpk + 190 days) V-band observation was obtained with the Inamori-Magellan Areal Camera and Spectrograph (IMACS, Dressler et al. 2006) mounted on the Magellan/Baade 6.5-m telescope on 2013 Apr 11.40. Using standard tasks in IRAF to perform aperture photometry and calibrating to a standard star field at similar airmass, we measure V = 19.65 ± 0.02 mag (exposure time of 90 s).
2.4. Optical spectroscopy
We obtained 28 epochs of optical spectroscopy of SN 2009ip covering the time period 2012 August 26 to 2013 April 11 using a number of facilities (see Table 3).
SN 2009ip was observed with the MagE (Magellan Echellette) Spectrograph mounted on the 6.5-m Magellan/Clay Telescope at Las Campanas Observatory. Data reduction was performed using a combination of Jack Baldwin’s mtools package and IRAF37 echelle tasks, as described in Massey et al. (2012). Optical spectra were obtained at the F. L. Whipple Observatory (FLWO) 1.5-m Tillinghast telescope on several epochs using the FAST spectrograph (Fabricant
37 IRAF is distributed by the National Optical Astronomy Observatories, which are operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.

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Flux (10−14 erg cm−2 s−1 Å−1)

1.2 C III]
1 Si III]

[N II]

Mg II

0.8 Fe II Fe II
0.6

0.4

0.2

10800 0.25

2000

2200 2400 2600 2800 Rest wavelength (Å)

3000

3200

Si IV Lyα Ti III/Fe II? 0.2
C II

Fe II+Al II

days). The spectra were simultaneously observed in three different arms, covering the entire wavelength range 3000–
25000 Å: ultra-violet and blue (UVB), visual (VIS) and nearinfrared (NIR) wavebands. The main dispersion was achieved through a 180 grooves/mm echelle grating blazed at 41.77◦ (UVB), 99 grooves/mm echelle grating blazed at 41.77◦ (VIS) and 55 grooves/mm echelle grating blazed at 47.07◦ (NIR). Observations were performed at parallactic angle under the following conditions: clear sky, the average seeing was ∼ 0.7 and ∼ 1.0 and the airmass range was ∼ 1.1–1.23. We used the X-shooter pipeline (Modigliani et al. 2010) in physical mode to reduce both SN 2009ip and the standard star spectra to two-dimensional bias-subtracted, flat-field corrected, order rectified and wavelength calibrated spectra in counts. To obtain 1-D spectra the 2-D spectra from the pipeline were optimally extracted (see Horne 1986) using a custom IDL program. Furthermore, the spectra were slit-loss corrected, flux calibrated and corrected for heliocentric velocities using a custom IDL program. The spectra were not (carefully) telluric corrected. The sequence of optical spectra is shown in Fig. 6.

Flux (10−14 erg cm−2 s−1 Å−1)

0.15 0.1

NV OI

Si II C IV

CI

0.05

0 1200

1300 1400 1500 1600 Rest wavelength (Å)

1700

FIG. 5.— Upper panel: HST-STIS spectrum obtained on 2012 October 29
(tpk + 26 days). The C III] and Si III] identifications are in a noisy part of the spectrum and are therefore uncertain. Lower panel: HST-COS spectrum
obtained on 2012 November 6 (tpk + 34 days).

et al. 1998). Data were reduced using a combination of standard IRAF and custom IDL procedures (Matheson et al. 2005). Low and medium resolution spectroscopy was obtained with the Robert Stobie Spectrograph mounted on the Southern African Large Telescope (SALT/RSS) at the South African Astronomical Observatory (SAAO) in Sutherland, South Africa. Additional spectroscopy was acquired with the Goodman High Throughput Spectrograph (GHTS) on the SOAR telescope. We also used IMACS mounted on the Magellan telescope and the Low Dispersion Survey Spectrograph 3 (LDSS3) on the Clay telescope (Magellan II). The Multiple Mirror Telescope (MMT) equipped with the “Blue Channel” spectrograph (Schmidt et al. 1989) was used to monitor the spectral evolution of SN 2009ip over several epochs. Further optical spectroscopy was obtained with the R-C CCD Spectrograph (RCSpec) mounted on the Mayall 4.0 m telescope, a Kitt Peak National Observatory (KPNO) facility. Spectra were extracted and calibrated following standard procedures using IRAF routines.
We used the X-shooter echelle spectrograph (D’Odorico et al. 2006) mounted at the Cassegrain focus of the Kueyen unit of the Very Large Telescope (VLT) at the European Southern Observatory (ESO) on Cerro Paranal (Chile) to obtain broad band, high-resolution spectroscopy of SN 2009ip on 2012 September 30 (tpk − 3 days) and October 31 (tpk + 28

2.5. NIR photometry
We obtained ZYJHK data using the Wide-field Infrared Camera (WFCAM) on the United Kingdom Infrared Telescope (UKIRT). The observation started on 2012 September 23 (tpk − 10 days), and continued on a nearly daily basis until 2012 December 31(tpk + 89 days) when SN 2009ip settled behind the Sun. The data reduction was done through an automatic pipeline of the Cambridge Astronomy Survey Unit. The flux of the object was measured with AUTO-MAG of SExtractor (Bertin & Arnouts 1996), where the photometric calibration was done using 2MASS stars within a radius of 8 arcmin from SN 2009ip. The 2MASS magnitudes of the stars were converted to the UKIRT system following Hodgkin et al. (2009) and the stars with the magnitude errors smaller than 0.10 mag were used for the photometry calibration, which yields typically 20-30 stars. A more detailed description of the NIR photometry can be found in Im et al. (in preparation).
Additional NIR photometry was obtained with PAIRITEL, the f/13.5 1.3-meter Peters Automated Infrared Imaging TELescope at the Fred Lawrence Whipple Observatory (FLWO) on Mount Hopkins, Arizona (Bloom et al. 2006). PAIRITEL data were processed with a single mosaicking pipeline that co-adds and registers PAIRITEL raw images into mosaics (see Wood-Vasey et al. 2008; Friedman 2012).Aperture photometry with a 3 aperture was performed at the SN position in the mosaicked images using the IDL routine aper.pro. No aperture corrections or host galaxy subtraction were performed. Figure 2 presents the complete SN 2009ip NIR data set. The PAIRITEL photometry can be found in Table 7. A table will the UKIRT photometry will be published in Im et al. (in preparation).
2.6. NIR spectroscopy
In addition to the X-shooter spectra, early time, lowdispersion (R ≈ 700) NIR spectra covering 0.9 to 2.4 µm were obtained with the 2.4m Hiltner telescope at MDM Observatory on 2012 September 27 (tpk − 6 days) and September 29 (tpk − 4 days). The data were collected using TIFKAM, a high-throughput infrared imager and spectrograph with a 1024 × 1024 Rockwell HgCdTe (HAWAII-1R) detector. The

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FIG. 6.— Optical spectra of SN 2009ip. Shades of red (blue) are used for spectra obtained during the rise time (decay time) of the 2012b explosion. Black is used for the 2012a outburst.

8 Margutti et al.
FIG. 7.— NIR spectral sequence with line identifications overlaid. The high-resolution spectrum obtained on 2012 November 19 has been smoothed here for display purposes. Shades of red (blue) have been used for spectra obtained during the rise (decay) time. A portion of the VLT/X-shooter spectra is also shown here.

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9

FIG. 8.— X-ray (Swift-XRT and XMM-Newton, filled and open circle, respectively) and 9 GHz radio light-curve (red squares, VLA) of SN 2009ip. X-rays are detected when the bolometric luminosity reaches its peak. Radio emission is detected at much later times. A re-scaled version of the bolometric light-curve is also shown for comparison. This plot does not include the late time X-ray limit obtained on 2013 April 4.5 (tpk + 183 days).
target was dithered along the 0.6 slit in a ABBA pattern to minimize the effect of detector defects and provide first-order background subtractions. Data reduction followed standard procedures using the IRAF software. Wavelength calibration of the spectra was achieved by observing argon lamps at each position. The spectra were corrected for telluric absorption by observing A0V stars at similar airmasses, and stellar features were removed from the spectra by dividing by an atmospheric model of Vega (Kurucz 1993).
Additional NIR low-resolution spectroscopy of SN 2009ip was obtained with the Folded-Port Infrared Echellette (FIRE) spectrograph (Simcoe et al. 2013) on the 6.5-m Magellan Baade Telescope, with simultaneous coverage from 0.82 to 2.51 µm. Spectra were acquired on 2012 November 5, 25 and December 3. The object was nodded along the slit using the ABBA pattern. The slit width was 0.6 , yielding R ≈ 500 in the J band. Data were reduced following the standard procedures described in Vacca et al. (2003), Foley et al. (2012) and Hsiao et al. (2013). An A0V star was observed for telluric corrections. The resulting telluric correction spectrum was also used for the absolute flux calibration.
Moderate-resolution (R ∼ 6000) NIR spectroscopy was obtained on 2012 November 19 with FIRE. SN 2009ip was observed in high-resolution echellette mode with the 0.6 slit. Eight frames were taken on source with 150 s exposures using ABBA nodding. Data were reduced using a custom-developed IDL pipeline (FIREHOSE), evolved from the MASE suite used for optical echelle reduction (Bochanski et al. 2009). Standard procedures were followed to apply telluric corrections and relative flux calibrations as described above. Finally, the corrected echelle orders were combined into single 1D spectrum for analysis. The complete sequence of NIR spectra is shown in Fig. 7. The observing log can be found in Table 4.
2.7. Millimeter and Radio Observations: CARMA and EVLA
We obtained two sets of millimeter observations at mean frequency of ∼84.5 GHz (∼7.5 GHz bandwidth) with the Combined Array for Research in Millimeter Astronomy (CARMA; Bock et al. 2006) around maximum light, begin-

TABLE 1 RADIO AND MILLIMETER OBSERVATIONS OF SN 2009IP

Date (UT start time)

Fν (µJy)

ν Instrument (GHz)

2012 Sep 2012 Sep 2012 Sepa 2012 Sep 2012 Oct 2012 Oct 2012 Oct 2012 Oct 2012 Nov 2012 Nov 2012 Dec 2012 Dec 2013 Mar

26.096 26.096 26.63 27.170 16.049 17.109 17.120 26.036
6.078 12.966 1.987 2.932
9.708

< 115.2 < 46.5
< 66 < 1000 < 70.5 < 104.1 < 1500 < 36.3 < 59.1 72.6 ± 15.2 < 70.5 78.3 ± 21.4
< 9.6

21.25 8.85 18 84.5
21.25 21.25
84.5 8.85 21.25 8.85 21.25 8.99 9.00

VLA VLA ATCA CARMA VLA VLA CARMA VLA VLA VLA VLA VLA VLA

NOTE. — Errors are 1σ and upper limits are 3σ.
aFrom Hancock et al. (2012).
ning 2012 September 27.17 (tpk − 6 days) and 2012 October 17.12 (tpk + 14 days). We utilized 2158-150 and 2258-279 for gain calibration, 2232+117 for bandpass calibration, and Neptune for flux calibration. In ∼160 and ∼120 min integration time at the position of SN 2009ip, we obtain 3-σ upper limits on the flux density of 1.5 and 1.0 mJy, respectively. The overall flux uncertainty with CARMA is ∼20%.
We observed the position of SN 2009ip with the Karl G. Jansky Very Large Array (VLA; Perley et al. 2011) on multiple epochs beginning 2012 September 26.10 (tpk − 7 days), with the last epoch beginning 2012 Dec 2.93 (tpk + 61 days). These observations were carried out at 21.25 GHz and 8.85 GHz with 2 GHz bandwidth in the VLA’s most extended configuration (A; maximum baseline length ∼36.4 km) except for the first observations, which were obtained in the BnA configuration. In most epochs our observations of flux calibrator, 3C48, were too contaminated with Radio Frequency Interference (RFI). Therefore, upon determining the flux of our gain calibrator J2213-2529 from the best observations of 3C48, we set the flux density of J2213-2529 in every epoch to be 0.65 and 0.63 Jy, for 21.25 and 8.9 GHz, respectively. We note that this assumption might lead to slightly larger absolute flux uncertainties than usual (∼15-20%). In addition, the source-phase calibrator cycle time (∼6 min) was a bit longer than standard for high frequency observations in an extended configuration, potentially increasing decoherence. We manually inspected the data and flagged edge channels and RFI, effectively reducing the bandwidth by ∼15%. We reduced all data using standard procedures in the Astronomical Image Processing System (AIPS; Greisen 2003). A summary of the observations is presented in Table 1.
No source is detected at the position of SN 2009ip at either frequencies during the first 50 days since the onset of the major outburst in September 2012, enabling deep limits on the radio emission around optical maximum. A source is detected at 8.85 GHz on 2012 November 13 (tpk +41 days), indicating a re-brightening of SN 2009ip radio emission at the level of Fν ∼ 70 µJy. The source position is α=22:23:08.29 ±0.01 and δ= −28:56:52.4 ±0.1 , consistent with the position determined from HST data. We merged the two observations that yielded a detection to improve the signal to noise and constrain the spectrum. Splitting the data into two 1 GHz slices centered at 8.43 GHz and 9.43 GHz, we find inte-

10 Margutti et al.

FIG. 9.— X-ray spectra of SN 2009ip. The Swift-XRT spectrum collects

observations obtained around the optical peak (tpk − 2 days until tpk + 13 days, total exposure time of 86 ks). The XMM EPIC-PN spectrum was obtained on

2012 November 3 (tpk + 31 days, total exposure of 55 ks). The spectral model

consists of absorbed bremsstrahlung emission at kT = 60 keV and intrinsic

athbesoXrpMtiMona(nNdHXinRtT=s0p.e1c0tr+−a00,..00r65e×spe1c0t2iv2eclmy)−.2

and NHint < 3.1 Contamination by

× 1021cm−2 for a nearby source

lying ≈ 6 from SN 2009ip is expected for the XMM spectrum. The color-

coded shaded areas highlight the presence of possible emission excess with

respect to the model.

FIG. 10.— Left panel: Swift-XRT image of the field of SN 2009ip collecting data before and after the optical peak (−32 days < t − tpk < −2 days and +29 days < t − tpk < +83 days), for a total exposure time of 110 ks. Right panel: same field imaged around the optical peak (−2 days < t − tpk < +13 days) for a total exposure of 86 ks. In both panels a white circle marks a 10 region around SN 2009ip. An X-ray source is detected at a position consistent with SN 2009ip around the optical peak (right panel) with significance of 6.1σ. The contaminating source discussed in the text is apparent in the left image.
grated flux densities of Fν = 60.0 ± 16.7 µJy (8.43 GHz) and Fν = 100.6 ± 18.9 µJy (9.43 GHz), suggesting an optically thick spectrum. The upper limit of Fν < 70.5 µJy at 21.25 GHz on 2012 December 2 indicates that the observed spectral peak frequency νpk is between 9.43 GHz and 21.25 GHz. A late-time observation obtained on March 9th shows that the radio source faded to Fν < 9.6 µJy at 9 GHz, pointing to a direct association with SN 2009ip. The radio light-curve at 9 GHz is shown in Figure 8.
2.8. X-ray observations: Swift-XRT and XMM-Newton
We observed SN 2009ip with the Swift/X-Ray Telescope (XRT, Burrows et al. 2005) from 2012 September 4

(20:36:23) until 2013 January 1 (13:43:55), for a total exposure of 260 ks, covering the time period −29 days < t − tpk < +90 days. Data have been entirely acquired in Photon Counting (PC) mode38 and analyzed using the latest HEASOFT (v6.12) release, with standard filtering and screening criteria. No X-ray source is detected at the position of SN 2009ip during the decay of the 2012a outburst (t < −11 days), down to a 3σ limit of 3 × 10−3cps in the 0.3-10 keV energy band (total exposure of 12.2 ks). Observations sampling the rise time of the 2012b explosion ( −11 days < t − tpk < −2 days ) also show no detection. With 31.4 ks of total exposure the 0.3-10 keV count-rate limit at the SN position is < 1.1 × 10−3cps. Correcting for PSF (Point Spread Function) losses and vignetting and merging the two time intervals we find no evidence for Xray emission originating from SN 2009ip in the time interval −29 days < t − tpk < −2 days down to a limit of < 5.6 × 10−4cps (0.3-10 keV, exposure time of 43.6 ks). X-ray emission is detected at a position consistent with SN 2009ip starting from tpk − 2 days, when the 2012b explosion approached its peak luminosity in the UV/optical bands (Fig. 8). The source is detected at the level of 5σ and 4σ in the time intervals −2 days < t − tpk < +3 days and +4 days < t − tpk < +13 days, respectively, with PSF and vignetting corrected count-rates of (1.6 ± 0.3) × 10−3cps and (6.8 ± 1.8) × 10−4cps (0.3-10 keV, exposure time of 42 and 44 ks).
Starting from tpk + 17 days, the source is no longer detected by XRT. We therefore activated our XMM-Newton program (PI P. Chandra) to follow the fading of the source. We carried out XMM-Newton observations starting from 2012 November 3 at 13:25:33 (tpk + 31 days). Observations have been obtained with the EPIC-PN and EPIC-MOS cameras in full frame with thin filter mode. The total exposures for the EPIC-MOS1 and EPIC-MOS2 are 62.62 ks and 62.64 ks, respectively, and for the EPIC-PN, the exposure time is 54.82 ks. A point-like source is detected at the position of SN 2009ip with significance of 4.5 σ (for EPIC-PN), and rate of (2.7 ± 0.3) × 10−3 cps in a region of 10 around the optical position of SN 2009ip.
From January until April 2013 the source was Sun constrained for Swift. 10 ks of Swift-XRT data obtained on 2013 April 4.5 (tpk + 183 days, when SN 2009ip became observable again) showed no detectable X-ray emission at the position of the transient down to a 3σ limit of 4.1 × 10−3 cps (0.3-10 keV).
We use the EPIC-PN observation to constrain the spectral parameters of the source. We extract photons from a region of 10 radius to avoid contamination from a nearby source (Fig. 10). The XMM-Newton software SAS is used to extract the spectrum. Our spectrum contains a total of 132 photons. We model the spectrum with an absorption component (which combines the contribution from the Galaxy and from SN 2009ip local environment, tbabs × ztbabs within Xspec) and an emission component. Both thermal bremsstrahlung and thermal emission from an optically thin plasma in collisional equilibrium (Xspec MEKAL model) can adequately fit the observed spectrum. In both cases we find kT > 10 keV and intrinsic hydrogen absorption of NHint ≈ 1021 cm−2. In the following we assume thermal bremsstrahlung emission with kT = 60 keV (this is the typical energy of photons expected from shock break-out from a dense CSM shell, see
38 The Swift-XRT observing modes are defined in Hill et al. (2004).

SN2009ip

11

Section 7.5).39 The Galactic absorption in the direction of SN 2009ip is NH,MW = 1.2 × 1020 cm−2 (Kalberla et al. 2005). The best-fitting neutral hydrogen intrinsic absorption40 is constrained to be NHint = 0.10+−00..0065 × 1022cm−2. Using these parameters, the corresponding unabsorbed (absorbed) flux is (1.9 ± 0.2) × 10−14erg s−1cm−2 ((1.7 ± 0.2) × 10−14erg s−1cm−2) in the 0.3 − 10 keV band. The spectrum is displayed in Fig. 9 and shows some evidence for an excess of emission around ∼ 7−8 keV (rest-frame) which might be linked to the presence of Ni or Fe emission lines (see e.g. SN2006jd and SN2010jl; Chandra et al. 2012a,b).
A Swift-XRT spectrum extracted around the peak (−2 days < t − tpk < +13 days, total exposure of 86 ks) can be fit by a thermal bremsstrahlung model, assuming kT = 60 keV and NHint < 3.1 × 1021cm−2 at the 3σ c.l. As for XMM, we use a 10 extraction region to avoid contamination from a nearby source (Fig. 10). The count-to flux conversion factor deduced from this spectrum is 3.8×10−11erg s−1cm−2ct−1 (0.310 keV, unabsorbed). We use this factor to calibrate our SwiftXRT light-curve. The complete X-ray light-curve is shown in Fig. 8.
We note that at the resolution of XMM and Swift-XRT we cannot exclude the presence of contaminating X-ray sources at a distance 10 . We further investigate this issue constraining the level of the contaminating flux by merging the Swift-XRT time intervals that yielded a non-detection at the SN 2009ip position. Using data collected between tpk − 29 days and tpk − 2 days, complemented by observations taken between tpk + 29 days and tpk + 90 days, we find evidence for an X-ray source located at RA=22h23m09.19s and Dec=−28◦56 48.7 (J2000), with an uncertainty of 3.8 radius (90% containment), corresponding to 1 from SN 2009ip. The source is detected at the level of 3.4σ with a PSF, vignetting and exposure corrected count-rate of (3.0 ± 0.9) × 10−4 cps (total exposure of 110 ks, 0.3-10 keV energy band). The field is represented in Fig. 10, left panel. This source contaminates the reported SN 2009ip flux at the level of ∼ 1.6 × 10−4 cps. Adopting the count-toflux conversion factor above, this translates into a contaminating unabsorbed flux of ∼ 6 × 10−15erg s−1cm−2 (luminosity of ∼ 5 × 1038erg s−1 at the distance of SN 2009ip), representing ∼ 10% the X-ray luminosity of SN 2009ip at peak. This source does not dominate the X-ray energy release around the peak time.
Observations obtained with the Chandra X-ray Observatory (PI D. Pooley) on tpk + 19 days reveal the presence of an additional X-ray source lying ≈ 6” from SN 2009ip and brighter than SN 2009ip at that time. SN 2009ip is also detected (Pooley, private communication). Our contemporaneous Swift-XRT observations constrain the luminosity of the contaminating source to be 1.5 × 1039erg s−1, 30% the X-ray luminosity of SN 2009ip at peak. We conclude that the contaminating source is not dominating the X-ray emission of SN 2009ip around peak, if stable. The temporal coincidence of the peaks of the X-ray and optical emission of SN 2009ip is suggestive that the detected X-ray emission is physically associated with SN 2009ip. However, given the
39 We consider a non-thermal power-law emission model unlikely given the very hard best-fitting photon index of Γ = 0.87 ± 0.15 we obtain from this spectrum.
40 This estimate assumes an absorbing medium with solar abundance and low level of ionization.

uncertain contamination, in the following we conservatively assume Lx 2.5 × 1039erg s−1 for the peak X-ray luminosity of SN 2009ip.
2.9. Hard X-ray observations: Swift-BAT
Stellar explosions embedded in an optically thick medium have been shown to produce a collisionless shock when the shock breaks out from the progenitor environment, generating photons with a typical energy 60 keV (Murase et al. 2011, Katz et al. 2011). We constrain the hard X-ray emission from SN 2009ip exploiting our Swift-BAT (Burst Alert Telescope, Barthelmy et al. 2005) campaign with observations obtained between 2012 September 4 (tpk − 29 days) and 2013 January 1 (tpk + 90 days) in survey mode (15-150 keV energy range). We analyzed the Swift-BAT survey data following standard procedures: sky images and source rates were obtained by processing the data with the BATSURVEY tool adopting standard screening and weighting based on the position of SN 2009ip. Following the BAT survey mode guidelines, fluxes were derived in the four standard energy channels, 14–24, 24–50, 50–100, and 100–195 keV. We converted the source rates to energy fluxes assuming a typical conversion factor of (5.9 ± 1.0) × 10−7 erg cm−2 s−1/count s−1, estimated assuming a range of different photon indices of a power–law spectrum (Γ = 1−3). In particular, analyzing the data acquired around the optical peak we find evidence for a marginal detection at the level of 3.5σ in the time interval −0.8 days < t − tpk < +0.2 days (corresponding to 2012 October 2.2 – 3.2). A spectrum extracted in this time interval can be fit by a power– law spectrum with photon index Γ = 1.8 ± 1.0 (90% c.l.) leading to a flux of (2.6 ± 1.4) × 10−10erg s−1cm−2 (90% c.l., 15150 keV, total exposure time of 7.0 ks). The simultaneity of the hard X-ray emission with the optical peak is intriguing. However, given the limited significance of the detection and the known presence of a non-Gaussian tail in the BAT noise fluctuations (H. Krimm, private communication), we conservatively use F < 7 × 10−10erg s−1cm−2 (L < 8 × 1040erg s−1) as the 5σ upper limit to the hard X-ray emission from SN 2009ip around maximum light, as derived from the spectrum above.
2.10. GeV observations: Fermi-LAT
GeV photons are expected to arise when the SN shock collides with a dense circumstellar shell of material, almost simultaneous with the optical light-curve peak (Murase et al. 2011, Katz et al. 2011). We searched for high-energy γ-ray emission from SN 2009ip using all-sky survey data from the Fermi Large Area Telescope (LAT; Atwood et al. 2009), starting from 2012 September 3 (tpk − 30 days) until 2012 October 31 (tpk + 28 days). We use events between 100 MeV and 10 GeV from the P7SOURCE_V6 data class (Ackermann et al. 2012), which is well suited for point-source analysis. Contamination from γ-rays produced by cosmic-ray interactions with the Earth’s atmosphere is reduced by selecting events arriving at LAT within 100° of the zenith. Each interval is analyzed using a Region Of Interest (ROI) of 12° radius centered on the position of the source. In each time window, we performed a spectral analysis using the unbinned maximum likelihood algorithm gtlike. The background is modeled with two templates for diffuse γ-ray background emission: a Galactic component produced by the interaction of cosmic rays with the gas and interstellar radiation fields of the Milky Way, and an isotropic component that includes both the contribution of the extragalactic diffuse emission and

12 Margutti et al.

the residual charged-particle backgrounds.41 We fix the normalization of the Galactic component but leave the normalization of the isotropic background as a free parameter. We also include the bright source 2FGL J2158.8−3013, located at approximately 5.◦48 from the location of SN 2009ip, and we fixed its parameters according to the values reported in Nolan et al. (2012).
We find no significant emission at the position of SN 2009ip. Assuming a simple power-law spectrum with photon index Γ = 2, the typical flux upper limits in 1-day intervals are (1 − 3) × 10−10 ergs cm−2 s−1 (100 MeV – 10 GeV energy range, 95% c.l.). Integrating around the time of the optical peak (−2 days < t − tpk < +4 days ) we find F < [2.1, 1.9, 3.6]×10−11 ergs cm−2 s−1 for three energy bands (100 MeV–464 MeV, 464 MeV–2.1 GeV and 2.1 GeV–10 GeV).42
3. EVOLUTION OF THE CONTINUUM FROM THE UV TO THE NIR
Our 13-filter photometry allows us to constrain the evolution of the spectral energy distribution (SED) of SN 2009ip with high accuracy. We fit a total of 84 SEDs, using data spanning from the UV to the NIR.
The extremely blue colors and color evolution of SN 2009ip (see Fig. 11, lower panel, and Fig. 12) impose non-negligible deviations from the standard UVOT count-to-flux conversion factors. The filter passbands (e.g. the presence of the "red leak" in the w2 and w1 filters) also affects the energy distribution of the detected photons for different incoming spectra. Because of the rapidly changing spectral shape in the UV, even the ratio of intrinsic flux to observed counts through the m2 filter, which has no significant red leak, is strongly dependent on the spectral shape. We account for these effects as follows: first, for each filter, we determine a grid of countto-flux conversion factors at the effective UVOT filters Vega wavelengths listed in Poole et al. 2008, following the prescriptions by Brown et al. (2010). We assume a black-body spectrum as indicated by our analysis of the SED of SN 2009ip at optical wavelengths. Our grid spans the temperature range between 2000 K and 38000 K with intervals of 200 K. We observe a variation in the conversion factor of 90%, 17%, 7% and 5% in the w2, m2 , w1 and u filters as the temperature goes from 6000 K to 20000 K. For the v and b filters the variation is below 1%. As a second step we iteratively fit each SED consisting of UVOT plus ground-based observations until the input black-body temperature assumed to calibrate the UVOT filters matches the best-fitting temperature within uncertainties.
For t > tpk − 7 days the UV+BVRI SED is well fitted by a black-body spectrum with a progressively larger radius ("hot" black-body component in Fig. 11). The temperature evolution tracks the bolometric luminosity, with the photosphere becoming appreciably hotter in correspondence with lightcurve bumps and then cooling down after the peak occurred. Around tpk + 70 days the temperature settles to a floor around 5000 K and remains nearly constant in the following 20 days. The temperature has been observed to plateau at similar times in some SNe IIn (e.g. SN 2005gj and SN 1998S where the black-body temperature reached a floor at ∼ 6000 − 6500 K;
41 The models used for this analysis, gal_2yearp7v6_v0.fits and iso_p7v6source.txt, are available from the Fermi Science Support Center, http://fermi.gsfc.nasa.gov/ssc/. This analysis uses the Fermi-LAT Science Tools, v. 09-28-00.
42 We note the presence of a data gap between tpk − 9 days and tpk − 2 days due to target-of-opportunity observations by Fermi during that time.

Prieto et al. 2007, Fassia et al. 2000) and in SNe IIP as well (e.g. SN 1999em, with a plateau at ∼ 5000 K; Leonard et al. 2002).
The black body radius increases from ∼ 5.1 × 1014 cm to ∼ 6.3 × 1014cm in ∼ 0.3 days (from t = tpk − 6.8 days to t = tpk − 6.5 days), then makes a transition to a linear evolution with average velocity of ∼ 4000 − 4500 km s−1 until tpk + 20 days, followed by a plateau around RHOT = 1.6 × 1015cm. In the context of the interaction scenario of Section 8 this change in the black-body radius evolution with time likely marks the transition to when the interaction shell starts to become optically thin (the black-body radius is a measure of the effective radius of emission: the shock radius obviously keeps increasing with time). A rapid decrease in radius is observed around tpk + 70 days. After this time the RHOT mimics the temporal evolution of the bolometric light-curve (see Fig. 11). In SNe dominated by interaction with pre-existing material, the black-body radius typically increases steadily with time, reaches a peak and then smoothly transitions to a decrease (see e.g. SN 1998S, Fassia et al. 2000; SN 2005gj, Prieto et al. 2007). The more complex behavior we observe for SN 2009ip likely results from a more complex structure of the immediate progenitor environment (Section 8).
Starting from tpk + 16 days, the best-fitting black body model tends to over-predict the observed flux in the UV, an effect likely due to increasing line-blanketing. As the temperature goes below ∼ 104 K, the recombination of the ejecta induces a progressive strengthening of metal-line blanketing which is responsible for partially blocking the UV light.We account for line-blanketing by restricting our fits to the UBVRI flux densities for t > tpk + 16 days. Our fits still indicate a rapidly decreasing temperature with time. We conclude that the rapid drop in UV light observed starting from tpk + 12 days mainly results from the cooling of the photosphere.
Starting around tpk + 59 days the UV emission fades more slowly and we observe a change in the evolution of the UV colors: from red to blue (Fig. 11, lower panel). The same evolution is observed in the (U-B) color of Fig. 12. This can also be seen from Fig. 2, where the NIR emission displays a more rapid decay than the UV. This manifests as an excess of UV emission with respect to the black-body fit.43 After tpk + 67 days a pure black-body spectral shape provides a poor representation of the UV to NIR SED.
We furthermore find clear evidence for excess of NIR emission with respect to the hot black body (see Fig. 13) as we first reported in Gall et al. (2012), based on the analysis of the VLT/X-shooter spectra (Fig. 19 and 20). Modeling the NIR excess with an additional black-body component, we obtain the radius and temperature evolution displayed in Fig. 11 ("cold" black body). The cold black-body radius is consistent with no evolution after tpk − 4 days, with RCOLD ∼ 4 × 1015cm. TCOLD is also found to be TCOLD ∼ 3000 K until tpk + 14 days, which implies LCOLD ≈ const for −4 days < t − tpk < +14 days (together with the almost unchanged NIR colors of Fig. 12).44 Starting from tpk +16 days TCOLD cools down to reach TCOLD ∼ 2000 K on tpk + 23 days. At this stage the hot black body with
43 This is especially true in the case of the UVOT m2 filter, which does not suffer from the "red leak".
44 This is also consistent with the almost flat K-band photometry. In this time interval the K-band photometry is dominated by the cold component. For t tpk + 12 days Lλ,HOT > Lλ,COLD at λ = λK, so that the K band flux starts to more closely follow the temporal evolution seen at bluer wavelengths.

SN2009ip

13

FIG. 11.— Upper panel: Bolometric light-curve of SN 2009ip calculated from the best-fitting black-body temperatures and radii displayed in the intermediate panels. Lower panel: UV color evolution with time. The onset of the 2012b explosion corresponds to a sudden change in UV colors. After that, the UV colors become progressively redder. In this plot, v and u refers to the optical photometry in the UVOT system. Vertical shaded bands mark the time of observed bumps in the photometry of Fig. 2: some are powerful enough to be clearly visible in the bolometric luminosity curve as well.

14 Margutti et al.

FIG. 12.— Upper panel: Optical colors. UVOT magnitudes have been converted into the Johnson filters using a dynamical correction that accounts for the evolution of the color of the source. Lower panel: NIR colors. While SN 2009ip clearly evolves towards redder optical colors starting from tpk − 3 days, no strong evolution is apparent in the NIR colors in the same time interval.
THOT ∼ 8500 K completely dominates the emission at NIR wavelengths and the fit is no longer able to constrain the parameters of the cold component. Our NIR spectra of Fig. 7 clearly rule out line-emission as a source of the NIR excess.
Applying the same analysis to the 2012a outburst we find that the temperature of the photosphere evolved from ∼ 13400 K (at tpk − 56 days) to 8000 K (tpk − 12 days), with an average decay of ∼ 120 K/day. Our modeling shows a slightly suppressed UV flux which we interpret as originating from metal line-blanketing. Notably, the SED at tpk − 38 days

FIG. 13.— Black solid line: best fitting SED model obtained at tpk −4.5 days which clearly shows the presence of the "hot" (red line) and "cold" (blue line) components in the spectrum.
(when we have almost contemporaneous coverage in the UBVRI and JHK bands) shows evidence for a NIR excess corresponding to TCOLD ∼ 2000 K at the radius consistent with RCOLD ∼ 4 × 1015cm (as found for the NIR excess during the 2012b explosion).
Finally, we use the SED best-fitting models above to compute the bolometric luminosity of SN 2009ip. Displayed in Fig. 11 is the contribution of the "hot" black body. The "cold" black-body contribution is marginal, being always (2 − 4)% the luminosity of the "hot" component.
4. SPECTRAL CHANGES AT UV/OPTICAL/NIR FREQUENCIES
Pastorello et al. (2012) find the spectrum of SN 2009ip during the 2012a outburst to be dominated by prominent Balmer lines. In particular, spectra collected in August and September 2012 show clear evidence for narrow emission components (FWHM≈ 800 km s−1 for Hα) accompanied by absorption features, indicating the presence of high velocity material with velocities extending to v ≈ −14000 km s−1 (Mauerhan et al. 2013). Our 2012 August 26 spectrum confirms these findings. SN 2009ip experienced a sudden re-brightening around 2012 September 23 (tpk − 10 days, Brimacombe 2012; Margutti et al. 2012), signaling the beginning of the 2012b explosion. By this time the Hα line developed a prominent broad emission component with FWHM≈ 8000 km s−1 (Mauerhan et al. 2013, their Fig. 5). The broad component disappeared 3 days later: our spectrum obtained on 2012 September 26 (tpk − 7 days) indicates that the Hα line evolved back to the narrow profile (Fig. 14), yet still retained evidence for absorption with a core velocity v ≈ −5000 km s−1, possibly extending to v ≈ −7500 km s−1. By 2012 September 30 (tpk − 3 days, Fig. 19 and 20) the spectrum no longer shows evidence for the high velocity components in absorption and is instead dominated by He I and H I lines with narrow profiles.
In the following months SN 2009ip progressively evolves from a typical SN IIn (or LBV-like) spectrum with clear signs of interaction with the medium, to a spectrum dominated by broad absorption features, more typical of SNe IIP (Fig. 21). Our two Xshooter spectra (Fig. 19 and 20) sample two key points in this metamorphosis, providing a broad band view of these spectral changes at high resolution. Broad features completely disappear by the time of our observations in April

SN2009ip

15

FIG. 14.— Evolution of the Hα line profile with time. Orange dashed line: emission components. Blue dot-dashed line: absorption components. Red thick line: composite line profile. The vertical blue lines mark the velocity of the absorption components.

2013 (tpk + 190 days, Fig. 31). At no epoch we find evidence for very narrow, low velocity blue shifted absorption at v ∼ −100 km s−1, differently from what typically observed in Type IIn SNe and LBVs (see e.g. SN 2010jl, Smith et al. 2012). The major spectral changes during the 2012b explosion can be summarized as follows:
• Broad/intermediate absorption/emission features progressively re-appear in the H Balmer lines, with evidence of multiple velocity components (Fig. 14, 15 and 16).
• Narrow He I lines weaken with time (Fig. 17); He I later re-emerges with the intermediate component only.
• Fe II features re-emerge and later develop P Cygni profiles.

in LBV-like eruptions, while being typical of a variety of SN explosions (Fig. 17). Broad absorption dips disappear ∼ 200 days after peak.
Around 100 days after peak, emission from forbidden transitions (see e.g. [CaII] λλ 7291, 7324 in Fig. 6) starts to emerge. At this time SN 2009ip settles behind the Sun. Despite limited spectral evolution between tpk + 100 days and tpk + 200 days (when SN 2009ip re-emerges from the Sun constraint) we do observe the absorption features to migrate to lower velocities. We discuss each of the items below.
Additional optical/NIR spectroscopy of SN 2009ip during the 2012b explosion has been published by Mauerhan et al. (2013), Pastorello et al. (2012), Levesque et al. (2012), Smith et al. (2013) and Fraser et al. (2013): we refer to these works for a complementary description of the spectral changes underwent by SN 2009ip.

• Emission originating from Na I D is detected (Fig. 17).
• A broad near-infrared Ca II triplet feature typical of Type IIP SNe develops starting around 2012 November 15 (Fig. 18).
• More importantly, SN 2009ip progressively develops broad absorption dips which have never been observed

4.1. Evolution of the H I line profiles
The Hα line profile experienced a dramatic change in morphology after the source suddenly re-brightened on 2012 September 23. Figure 14 shows the Hα line at representative epochs: at any epoch the Hα line has a complex profile resulting from the combination of a narrow (Lorentzian) component (FWHM< 1000 km s−1), intermediate/broad width

16 Margutti et al.
FIG. 15.— Hα, Hβ and Hγ line profiles of SN 2009ip at representative epochs. The blue dotted lines mark the velocity of the major absorption components identified by our fits of the Hα line profile. For 2012 December 5 we also added two absorption components at −1000 km s−1 and −4000 km s−1 identified in the Hβ and Hγ lines. The late-time spectrum acquired on tpk + 101 days (2013 January 12) shows limited evolution in the Hα profile with respect to the previous epoch and it is not shown here.

SN2009ip

17

FIG. 16.— Evolution of the Hα line with time. We model the Hα line with a combination of Lorentzian and Gaussian profiles. Orange markers represent the narrow Lorentzian profile. Red is used for the broad component, while blue is associated with the blue-shifted absorption components. Panel (a): bolometric light-curve for reference. Panel (b): the red (orange) bars span the FWHM of the broad (narrow) component. These values are also reported in panel (c). Blue dots: absorption minima as obtained by modeling the absorption with a combination of Gaussians. Negative values indicate blueshifted components. We use light-blue bars to mark the 1σ width as obtained from the fit. Panel (d) shows the evolution of the equivalent width of the narrow (orange), broad (red) and absorption (blue) components. The peak of the narrow component progressively shifts to larger redshifted velocities as illustrated in Panel (e) and independently found by Fraser et al. (2013). The vertical dashed line marks an important time in the evolution of SN 2009ip from different perspectives: from this plot it is clear that around this time the width of broad component undergoes a remarkable transition from FWHM∼ 3000 km s−1 to FWHM∼ 10000 km s−1.
components (FWHM> 1000 km s−1) and blue absorption features with evidence for clearly distinguished velocity components. Emission and absorption components with similar velocity are also found in the Hβ and Hγ line profiles (Fig. 15).
The evolution of the line profile results from changes in the relative strengths of the different components in addition to the appearance (or disappearance) of high-velocity blue absorption edges. The evolution of the width and relative

FIG. 17.— Key spectral changes in SN 2009ip between 4500−6200 Å: narrow He I emission lines subside while Na I D emission grows in strength. He I later re-appears with the broad/intermediate component. Fe II emission lines emerge, while broad absorption dips develop red-wards the Na I D (and
He I) lines, around 5650 Å. Starting from ∼ tpk + 30 days, additional broad
absorption features around 5770 Å and 5850 Å appear, associated with the He I and the Na I D lines; H I lines develop strong absorption on their blue wing.

18 Margutti et al.

FIG. 18.— Beginning at tpk + 30 days SN 2009ip develops broad emission
and absorption components between 8300 Å and 9000 Å we attribute to Ca II. NIR emission from the CaII triplet is typical of IIP SNe (e.g. Pastorello et al. 2006).
strength of the different components is schematically represented in Fig. 16. The broad component dominates over the narrow emission starting from tpk + 33 days and reaches its maximum width at tpk + 51 days. After this time, the width of the broad component decreases. There is evidence for an increasing width of the narrow component with time, accompanied by a progressive shift of the peak to higher velocities. Finally, high-velocity (v > 1000 km s−1) absorption features get stronger as the light-curve makes the transition from the rise to the decay phase. The spectral changes are detailed below.
By tpk − 7 days the broad components dominating the line profile 10 days before (Mauerhan et al. 2013) have weakened to the level that most of the emission originates from a much narrower component which is well described by a Lorentzian profile with FWHM≈ 1000 km s−1. Absorption from high velocity material (v ≈ −5000 km s−1, measured at the minimum of the absorption feature) is still detected when the 2012b explosion luminosity is still rising. The highresolution spectra collected on tpk − 5 and tpk − 4 days allow us to resolve different blue absorption components: modeling these absorption features with Gaussians, the central velocities are found to be v ≈ −2200 km s−1, ≈ −4000 km s−1, ≈ −5300 km s−1, ≈ −7500 km s−1 with σ ≈ 300 − 500 km s−1. These absorption features are detected in the Hβ and Hγ lines as well (Fig. 15). The width of the narrow component of emission decreases to FWHM≈ 280 km s−1.
On 2012 September 30 (tpk − 3 days) SN 2009ip approaches

its maximum luminosity (Fig. 11). From our high-resolution
spectrum the Hα line is well modeled by the combination of two Lorenztian profiles with FWHM≈ 240 km s−1 and FWHM≈ 2600 km s−1. We find no clear evidence for absorp-
tion components. Interpreting the broad wings as a result of
multiple Thomson scattering in the circumstellar shell of the
narrow-line radiation (Chugai 2001) suggests that the opti-
cal depth of the unaccelerated circumstellar shell envelope to Thomson scattering is τ ∼ 3.
High-velocity absorption features in the blue wing of
the Hα line progressively re-appear as the luminosity of
the explosion enters its declining phase. Eight days af-
ter peak the Hα line exhibits a combination of narrow (FWHM≈ 340 km s−1) and broad (FWHM≈ 2000 − 3000 km s−1) Lorentzian profiles and a weak P Cygni profile with an absorption minimum around −1600 km s−1. Three days later (tpk + 11 days) the broad component (FWHM≈ 2600 km s−1) of emission becomes more prominent while the
width of the narrow Lorentzian profile decreases again to FWHM≈ 220 km s−1. At this time the bolometric light-curve
exhibits a third bump (Fig. 11). High-velocity absorption
features re-appear in the blue wing of the Hα line with absorption minima at v ≈ −12000 km s−1 and v ≈ −8000 km s−1 (σ ∼ 1000 km s−1). The low velocity P Cygni absorption is also detected at v ≈ −1200 km s−1. The Hβ and Hγ lines
possibly show evidence for an additional absorption edge at v ≈ −4000 km s−1 (Fig. 15).
A lower resolution spectrum obtained on tpk + 18 days shows the development of an even stronger broad emission component with FWHM≈ 9400 km s−1. While we cannot re-
solve the different components of velocity responsible for
the blue absorption, we find clear evidence for a deep minimum at v ≈ −10000 km s−1 with edges extending to v ≈ −14000 − 15000 km s−1. The broad emission component keeps
growing with time: at tpk + 42 days it clearly dominates the Hα profile. At this epoch the Hα line consists of a narrow component with FWHM≈ 240 km s−1, a broad emission component (FWHM≈ 10600 km s−1) and a series of absorption
features on the blue wing (both at high and low velocity).
Our high-resolution spectrum resolve the absorption minima at v ≈ −12500 km s−1, ≈ −10000 km s−1, ≈ −7000 km s−1 and ≈ −1300 km s−1 (Fig. 14). The Hβ and Hγ lines exhibit an additional blue absorption at v ≈ 3000 km s−1 (Fig. 15).
By tpk + 51 days the broad component which dominates the Hα line reaches FWHM≈ 13600 km s−1. High velocity absorption features are still detected at v ≈ −12500 km s−1 and ≈ −10000 km s−1. The absorption feature at v ≈ −7000 km s−1
becomes considerably more pronounced and shows clear evidence for two velocity components with minima at v ≈ −7900 km s−1 and v ≈ −5700 km s−1. The low-velocity absorbing component is also detected with a minimum at v ≈ −1000 km s−1. A spectrum obtained 63 days after maximum shows little evolution in the Hα profile, the only difference being a more pronounced absorption at v ≈ −5700 km s−1.
At tpk + 79 days we find a less prominent broad component: by this time its width decreased from FWHM≈ 12500 km s−1 to FWHM≈ 8200 km s−1. A spectrum obtained at tpk + 101 days confirms this trend (FWHM of the broad component ≈ 7500 km s−1): the bulk of the absorption is now at lower velocities v ≈ −3100 km s−1 (with a tail possibly extending to v ≈ −8000 km s−1). At tpk + 190 days the blue-shifted absorption is

SN2009ip

19

FIG. 19.— High-resolution VLT/X-shooter spectra captured the evolution of SN 2009ip in fine detail. Upper panel: around the optical peak, on 2012 September 30 (tpk − 3 days), SN 2009ip shows a narrow-line dominated spectrum typical of SNe (and LBVs) interacting with a medium. One month later (lower panel of each plot) SN 2009ip started to develop broad emission components (see in particular the Hα line) and deep absorption features (e.g. the yellow-shaded band
around 5650 Å ) more typical of SNe IIP. The complementary 10000 − 24500 Å wavelength range is shown in Fig. 20. Data have been corrected for Galactic extinction.

20 Margutti et al.
FIG. 20.— High-resolution VLT/X-shooter spectra from 10000 to 24500 Å . Continued from Fig. 19.

2 Narrow-line dominated
1.5

SN2009ip

2009ip (-4 d) 1996L (Type IIn, +9 d)

21

Normalized Flux

1 Broad-line dominated
0.5

2009ip (+ 51 d) 2006bp (Type IIP, +73 d)

0 4000

5000

6000

7000

Rest Wavelength [Angstroms]

8000

FIG. 21.— SN 2009ip evolved from a narrow-line dominated spectrum typical of Type IIn SN explosions to a spectrum that clearly shows broad absorption features more typical of Type IIP SNe. Here we show the spectrum of Type IIn SN 1996L (Benetti et al. 1999) and Type IIP SN 2006bp (Quimby et al. 2007).

found peaking at even lower velocities of v −2400 km s−1, and the "broad" (now intermediate) component has FWHM of only ≈ 2000 km s−1.
Finally, comparing the H Paschen and Brackett emission lines using our two highest resolution spectra collected around the peak (narrow-line emission dominated spectrum at tpk − 3 days) and 28 days after peak (when broad components start to emerge, see Fig. 19 and 20), we find that for both epochs the line profiles are dominated by the narrow component (FWHM≈ 170 km s−1) with limited evolution between the two. The Paschen β line clearly develops an intermediatebroad component starting from tpk + 33 days (see Fig. 7). Spectra obtained by Pastorello et al. (2012) before the sudden re-brightening of 2012 September 23 (tpk − 10 days) show a similar narrow plus broad component structure, with the broad emission dominating the narrow lines between 2012 August 26 and 2012 September 23. As for the H Balmer lines, the broad component completely disappeared as the light-curve approached its maximum.
We conclude by noting that, observationally, tpk + 17 days (i.e. 2012 October 20) marks an important transition in the evolution of SN 2009ip: around this time the broad Hα component evolves from FWHM∼ 3000 km s−1 to FWHM∼ 10000 km s−1 (Fig. 16); the photospheric radius RHOT flattens to RHOT ∼ 1.6 × 1015cm while the hot black-body temperature transitions to a milder decay in time (Section 3, Fig. 11). It is intriguing to note that our modeling described in Section 7 independently suggests that this is roughly the time when the explosion shock reaches the edge of the dense shell of material previously ejected by the progenitor.
4.2. The evolution of He I lines

Conspicuous He I lines are not unambiguously detected in our spectrum obtained on 2012 August 26. They are, however, detected in our spectrum acquired one month later, ∼ 3 days after SN 2009ip re-brightened45. At this epoch the light curve of SN 2009ip is still rising. Similarly to H Balmer
lines, HeI features (the brightest being at 5876 Å , and 7065 Å46) exhibit a combination of a narrow-intermediate profile (FWHM≈ 1000 km s−1), a weak broad component (FWHM≈ 5000 km s−1) together with evidence for a P Cygni absorption at velocity v ≈ −5000 km s−1.
As for the H Balmer lines, high-resolution spectroscopy obtained at tpk − 5 and tpk − 4 days shows the appearance of multiple absorption components on the blue wing of the He I λ5876 and λ7065 lines, with velocities v ≈ −2000 km s−1, ≈ −4800 km s−1 and ≈ −7000 km s−1 measured at the absorption minima (to be compared with Fig. 15). High velocity absorption features disappear by tpk − 3 days: He I λ5876 and λ7065 show the combination of a narrow plus broader intermediate Lorentzian profiles with FWHM≈ 2000 km s−1 and FWHM≈ 240 km s−1, respectively.
Starting from tpk − 3 days, He I features become weaker until He I λ7065 is not detected in our high-resolution spectrum acquired at tpk + 28 days (Fig. 19 and Fig. 19). He I later reappears in our spectra taken in the second half of November (t > tpk + 43 days) showing the broad/intermediate component only (FWHM≈ 2500 km s−1 as measured at tpk + 63 days). At tpk + 79 days He I λ7065 shows an intermediate-broad emis-
45 Note that He I was clearly detected during the LBV-like eruption episodes in 2011 (Pastorello et al. 2012)
46 We also detect He I λ4713 (weak), He I λ5016 (later blended with Fe II λ5019), He I λ6678, on the red wing of Hα, He I λ7281 (weak) and He I λ10830 (blended with Paγ). He I λ5876 is also blended with Na I D emission.

22 Margutti et al.

4.4. The NIR Ca II feature
Starting from ∼ 30 days after peak, our spectra (Fig. 18) show the progressive emergence of broad NIR emission originating from the Ca II triplet λλ8498, 8542, 8662 (see also Fraser et al. 2013, their Fig. 4). The appearance of this feature is typically observed during the evolution of Type II SN explosions (see e.g. Pastorello et al. 2006). Interestingly, no previous outburst of SN 2009ip showed this feature (2012a outburst included, see Pastorello et al. 2012). No broad Ca II triplet feature has ever been observed in an LBV-like eruption.
Figure 18 also sjows the emergence of broad absorption
dips around 8400 Å and 8600 Å. If Ca II λ8662 is caus-
ing the absorption around 8600 Å, the corresponding velocity at the absorption minimum is v ≈ −2400 km s−1. This absorption developed between 51 days and 63 days after peak.
The absorption at λ ≈ 8400 Å is instead clearly detected in our spectra starting from tpk + 45 days and likely results from
the combination of OI and CaII. If OI (8447 Å) is dominating the absorption at minimum, the corresponding velocity is v ≈ −1500 km s−1.

FIG. 22.— Evolution of the broad absorption features associated with Hα, He I and Na I D lines. As the photosphere recedes into the ejecta broad absorption features appear in the spectra with 3 typical velocities: v ∼ −12000 km s−1 (red band); v ∼ −5500 km s−1 (orange band) and v ∼ −2500 km s−1 (yellow band). Absorption features at higher velocity are revealed at earlier times: we clearly detect material with v ∼ −12000 km s−1 starting around tpk + 9 days; material with v ∼ −5500 km s−1 starts to be detected around tpk + 28 days, while slowly moving ejecta with v ∼ −2500 km s−1 is only detected ∼ 60 days after peak. Left panel: velocity profile of the absorption arising from Na I D plus He I (λ = 5876 Å), compared with HeI 1.08 µm and HeI 2.06 µm velocity profiles. Right panel: Na I
D plus He I (λ = 5876 Å) vs. Hα velocity profile.
sion profile with FWHM≈ 3000 km s−1. A similar value is obtained at tpk + 101 days. Roughly 100 days later, on 2013
April 11 He I 7065 Å is clearly detected with considerably narrower emission (FWHM≈ 1000 km s−1). He I λ6678 also re-emerges on the red wing of the Hα profile (Fig. 31).
4.3. The evolution of Fe II lines
A number of Fe II lines from different multiplets have been observed during previous SN 2009ip outbursts (both in 2009, 2011 and the 2012a outburst, see Pastorello et al. 2012, their Fig. 5 and 6). The Fe responsible for this emission is therefore pre-existent the 2012 explosion. Fe II is instead not detected in our spectra until tpk + 17 days (Fig. 17). From the Xshooter spectrum acquired at tpk + 28 days we measure the FWHM of the narrow Fe II lines λ5018 and λ5169 (multiplet 42): FWHM≈ 240 km s−1. A similar value has been measured by Pastorello et al. (2012) from their 2012 August 18 and September 5 spectra. As a comparison, the FWHM of the narrow (Lorentzian) component of the Hα line measured from the same spectrum is ≈ 170 km s−1. By tpk + 63 days the Fe II emission lines develop a P Cygni profile (Fig. 17), with absorption minimum velocity of v ≈ −1000 km s−1, possibly extending to v ≈ −4000 km s−1.

4.5. The development of broad absorption features
High-velocity, broad absorption features appear in our spectra starting 9 days after peak (see yellow bands in Fig. 7, Fig. 17, Fig. 19, Fig. 20). Absorption features of similar strength and velocity have never been associated with an LBV-like eruption to date, and are more typical of SNe (Fig. 21). These absorption features are unique to the 2012b explosion and have not been observed during the previous outbursts of SN 2009ip (see Smith et al. 2010b, Foley et al. 2011, Pastorello et al. 2012).
As the photosphere recedes into the ejecta it illuminates material moving towards the observer with different velocities. Our observations identify He I, Na I D and H I absorbing at 3 typical velocities (Fig. 22). The blue absorption edge of He I plus Na I D extends to v ≈ 18000 km s−1, as noted by Mauerhan et al. (2013). High-velocity v ∼ −12000 km s−1 absorption appears first, around tpk + 9 days followed by the v ∼ −5500 km s−1 absorption around tpk + 28 days, which in turn is followed by slower material with v ∼ −2500 km s−1, seen in absorption only starting from ∼ tpk + 60 days. This happens since material with lower velocity naturally overtakes the photosphere at later times. Material moving at three distinct velocities argues against a continuous distribution in velocity of the ejecta and suggests instead the presence of distinct shells of ejecta expanding with typical velocity v ∼ −12000 km s−1, v ∼ −5500 km s−1 and v ∼ −2500 km s−1.
4.6. UV spectral properties
Our Swift-UVOT low-resolution spectroscopic monitoring campaign maps the evolution of SN 2009ip during the first month after its major peak in 2012 (Fig. 3). We do not find evidence for strong spectral evolution at UV wavelengths (Fig. 4): as time proceeds the Fe III absorption features become weaker while Fe II develops stronger absorption features, consistent with the progressive decrease of the blackbody temperature with time (Fig. 11). UVOT spectra show the progressive emergence of an emission feature around
2500 − 3000 Å that is later well resolved by HST/STIS as emission from Mg II λλ2796, 2803 lines as well as Fe II multi-
plets at ∼ 2550, 2630, 2880 Å (Fig. 5). The Mg II line profiles

SN2009ip

23

are similar to the H I line profiles, with a narrow component and broad, blue-shifted absorption features. As for the H I lines, the narrow component originates from the interaction with slowly moving CSM. We further identify strong, narrow emission from N II] at λλ2140, 2143. Emission from C III] (λ1909) and Si III] (λλ1892, 1896) might also be present, but the noise level does not allow a firm identification.
At shorter wavelengths, the HST/COS spectrum taken 34 days after peak shows a mixture of high and low ionization lines (Fig. 5, lower panel). We identify strong lines of C II (λλ1334.5, 1335.7), O I (λλ 1302.2-1306.0), Si II (λλ 1526.7, 1533.5). Of the higher ionization lines one notes C IV (λλ1548.2, 1550.8) and N V (λλ1238.8, 1242.8). Interestingly, N IV] λ1486.5 is either very weak or absent which indicates a medium with density n 109 cm−3. Fe II is also present, although the identification of the individual lines is
not straightforward (e.g. the Fe II feature at ∼ 1294 Å may also be consistent with Ti III). Lyα emission is also very well detected.
Around this time, both the optical, NIR and UV spectra are dominated by permitted transitions: in particular, despite the presence of high ionization lines there are no forbidden lines of, e.g., [O III] λλ 4959, 5007, N IV] λλ1486 or O III] λλ 1664, consistent with the picture of high density in the line forming region. (The [Ca II] λ 7300 lines will clearly emerge only after tpk + 79 days). The main exceptions are the [N II] λλ 2140, 2143 lines (Fig. 5). The explanation could be a comparatively high critical density, ∼ 3×109 cm−3 in combination with a high N abundance.
A comparison of high (C IV λλ1548.2, 1550.8 and N V λλ1238.8, 1242.8) and low (C II λ 1335) ionization emission line profiles in velocity space reveals no significant difference: the three lines extend to ∼ 850 km s−1 on the red side, while there is an indication of a somewhat smaller extent on the blue wing, ∼ 500 km s−1. This is however complicated by the P Cygni absorption features and the doublet nature of the C IV and N V lines. The mixture of low and high ionization lines indicates that there are several components present in the line emitting region. This may either be in the form of different density components, or different ionization zones. The similar line profiles argue for a similar location of the ionization zones, supporting the idea of a complex emission region with different density components. The observed X-ray emission can in principle be responsible for the ionization.
5. METALLICITY AT THE EXPLOSION SITE AND HOST ENVIRONMENT
The final fate of a massive star is controlled by the mass of its helium core (e.g. Woosley et al. 2007), which is strongly dependent on the initial stellar mass, rotation and composition. Metallicity has a key role in determining the mass-loss history of the progenitor, with low metallicity generally leading to a suppression of mass loss, therefore allowing lowermass stars to end their lives with massive cores. SN 2009ip is positioned in the outskirts of NGC 7259 (Fig. 1). The remote location of SN 2009ip has been discussed by Fraser et al. (2013). Our data reveal no evidence for an H II region in the vicinity of SN 2009ip that would allow us to directly measure the metallicity of the immediate environment. Thus, we inferred the explosion metallicity by measuring the host galaxy metallicity gradient. The longslit was placed along the galaxy center at parallactic angle. We extracted spectra of the galaxy at positions in a sequence across our slit, producing a set of integrated light spectra from ∼ 0 − 2 kpc from either side of

log(O/H)+12 (PP04N2)






 
	



SN 2009ip

      Galactocentric radius (kpc)
FIG. 23.— Metallicity profile of NGC 7259, the host galaxy of SN 2009ip, as derived from our long slit spectroscopy. The solid line shows the best fit metallicity gradient. The dashed horizontal line marks the solar metallicity and the vertical dashed line marks the SN galactocentric radius. The colors distinguish measurements from opposite sides of the galaxy center. The error bars reflect propagation of the emission line flux uncertainties only. This analysis constrains the metallicity at the explosion site of SN 2009ip to be 8.2 < log(O/H) + 12 < 8.6 (0.4 Z < Z < 0.9 Z ).
the galaxy center. We use the “PP04 N2” diagnostic of Pettini & Pagel
(2004) to estimate gas phase metallicity using the Hα and [N II] λ6584 emission lines. We estimate the uncertainty in the metallicity measurements by Monte Carlo propagation of the uncertainty in the individual line fluxes. The median uncertainty is 0.09 dex, which is similar to the systematic uncertainty in the calibration of the strong line diagnostic (Kewley & Ellison 2008). Robust metallicity profiles can not be recovered in other diagnostics due to the faintness of the [O III] lines in our spectroscopy.
Figure 23 shows the resulting metallicity profile of NGC 7259. The metallicity at the galaxy center is log(O/H) + 12 =∼ 8.8, ∼ 1.3 Z on the PP04 N2 scale, but declines sharply with radius. The metallicity profiles on each side of the galaxy center in our longslit spectrum are consistent. We therefore assume that the metallicity profile is azimuthally symmetric. We estimate the metallicity gradient by fitting a linear profile. The best fit gradient intercept and slope are 8.8 ± 0.02 dex and −0.11 ± 0.02 dex kpc−1, respectively.
SN 2009ip is located ∼ 43.4” from the center of the galaxy NGC 7259 (equal to ∼ 5.0 kpc at dL = 24 Mpc). This is more than twice the distance to which our metallicity profile observations extend. Extrapolating directly from this gradient would imply an explosion site metallicity of log(O/H) + 12 =∼ 8.2, or ∼ 0.4 Z . This metallicity would place SN 2009ip at the extreme low metallicity end of the distribution of observed host environments of Type II SN (Stoll et al. 2012), and nearer to the low metallicity regime of broad-lined Type Ic supernovae (Kelly & Kirshner 2012; Sanders et al. 2012). However, the metallicity properties of galaxies at distances well beyond a scale radius have not been well studied. It is likely that a simple extrapolation is not appropriate, and the metallicity profile in the outskirts of the galaxy may flatten

24 Margutti et al.

FIG. 24.— Absolute R-band magnitude of SN 2009ip (pink dots) compared to the sample of Type IIn SNe with R-band photometric coverage around the peak (from the literature). For SNe 1994W, 1998S, 1999el, 2003ma, 2005cp, 2005cl, 2005db, 2005gl and 2006gy we refer the reader to Kiewe et al. (2012) and references therein. The photometry of PTF09UJ has been presented by Ofek et al. (2010). Colored area: typical absolute magnitude of LBV-like eruption episodes. Pink triangles mark the luminosity of SN 2009ip during the 2011 outburst (t = 0 corresponds here to 2011 July 8 for convenience). The exceptional SN (impostor?) SN1961V (photographic plate magnitudes, from Pastorello et al. 2012) and η Carinae during the 19th century Great Eruption (Mvis as compiled by Frew 2004, see also Smith & Frew 2011) are shown with green triangles. The comparison with SN 2010mc is shown in Fig. 30.
(Werk et al. 2011) or drop significantly (Moran et al. 2012). In either case, it is unlikely that the explosion site metallicity is significantly enriched relative to the gas we observe at R ∼ 2 kpc, with log(O/H) + 12 ∼ 8.6 (∼ 0.9 Z ). If we adopt this value as the explosion site metallicity, it is fully consistent with the observed distribution of SNe II, Ib, and Ic (Kelly & Kirshner 2012; Sanders et al. 2012; Stoll et al. 2012).
Our best constraint on the explosion site metallicity is therefore 0.4 Z < Z < 0.9 Z , pointing to a (mildly) sub-solar environment.
6. ENERGETICS OF THE EXPLOSION
The extensive photometric coverage (both in wavelength and in time) gives us the opportunity to accurately constrain the bolometric luminosity and total energy radiated by SN 2009ip. SN 2009ip reaches a peak luminosity of Lpk = (1.7 ± 0.1) × 1043erg s−1 (Fig. 11). The total energy radiated during the 2012a outburst (from 2012 August 1 to September 23) is (1.5 ± 0.4) × 1048erg while for the 2012b explosion we measure Erad2 = (3.2 ± 0.3) × 1049 erg. As much as ∼ 35% of this energy was released before the peak, while 50% of Erad2 was radiated during the first ∼ 15 days. Subsequent re-brightenings (which constitute a peculiarity of SN 2009ip) only contributed to small fractions of the total energy.

FIG. 25.— Re-normalized R-band magnitude of SN 2009ip compared with the sample of SNe IIn of Fig. 24. Shades of orange (blue) are used for SNe with a slower (faster) rise time. SN 2009ip is characterized by a fast rise and fast decay.
The peak luminosity of SN 2009ip is not uncommon among the heterogeneous class of SNe IIn, corresponding to MR ≈ −18 mag (Fig. 24). Its radiated energy of (3.2 ± 0.3) × 1049 erg falls instead into the low energy tail of the distribution mainly because of the very rapid rise and decay times of the bolometric luminosity (Fig. 25). The limited energy radiated by SN 2009ip brings into question the final fate of the progenitor star: was the total energy released sufficient to fully unbind the star (i.e., terminal explosion) or does SN 2009ip results from a lower-energy ejection of only the outer stellar envelope (i.e., non-terminal explosion)? This topic is explored in Section 8. Indeed, stars might be able to survive eruptive/explosive events that reach a visual absolute magnitude of Mvis ≈ −17 mag (e.g. SN 1961V in Fig. 25, Van Dyk & Matheson 2012, Kochanek et al. 2011), so that the peak luminosity is not a reliable indicator of a terminal vs. non-terminal explosion.47 With an estimated radiated energy of 3.2 × 1049 erg (Davidson & Humphreys 1997) the "Great Eruption" of η-Carinae (see e.g. Smith 2013 and references therein) demonstrated that it is also possible to survive the release of comparable amount of energy, even if on time scales much longer than those observed for SN 2009ip (the "Great Eruption" lasted about 20 yrs).
SN 2010mc shows instead striking similarities to SN 2009ip both in terms of timescales and of released energy (Fig. 30). As in SN 2009ip, a precursor bump was detected ∼ 40 days before the major outburst. Ofek et al. (2013b) calculate Erad > 6 × 1047erg (precursor-bump) and Erad ∼ 3 × 1049erg (major outburst) for SN 2010mc, compared with Erad1 = (1.5± 0.4) × 1048erg and Erad2 = (3.2 ± 0.3) × 1049erg we calculated above for SN 2009ip. The very close similarity of SN 2010mc and SN 2009ip originally noted by Smith et al. (2013) has important implications on the nature of both explosions (see Section 8).
47 The same line of reasoning applies to the velocity of the fastest moving material measured from optical spectroscopy as pointed out by Pastorello et al. (2012): very fast material with v ∼ 12500 km s−1 was observed on the LBV-like outburst of SN 2009ip of September 2011, proving that highvelocity ejecta can be observed even without a terminal explosion.

SN2009ip

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7. SOURCE OF ENERGY AND PROPERTIES OF THE IMMEDIATE ENVIRONMENT
In the previous sections we concentrated on the properties of the explosion (e.g. energetics, evolution of the emission/absorption features) and of the environment (i.e. the metallicity) that can be directly measured; here we focus on properties that can be inferred from the data.
The light-curve of SN 2009ip shows two major episodes of emission: the precursor bump (2012a outburst) and the major re-brightening (2012b explosion). Is this phenomenology due to two distinct explosions or is the double-peaked light-curve the result of a single explosion? The main argument against a single explosion producing the two peaks is the observed evolution of the photospheric radius in Fig. 11. In the singleexplosion scenario material can only decelerate with time: at tpk + 7 days the photospheric radius is RHOT ∼ 1.2 × 1015cm and the velocity is v = 4500 km s−1. Extrapolating back in time, this implies that the zero-time of the 2012b explosion is later than tpk − 24 days. This is much later than the observed onset of the 2012a outburst that occurred at tpk − 56 days and favors against a single-explosion scenario. Models where the first bump in the light-curve is a SN explosion while the second peak is due to the interaction of the SN ejecta with the CSM (Mauerhan et al. 2013) belong to this category. In the following we proceed instead with a two-explosion hypothesis and argue that we witnessed two separate but causally connected explosions from the progenitor of SN 2009ip.
7.1. Limit on the Nickel mass synthesized by the 2012b explosion
Narrow emission lines in the optical spectra of SN 2009ip require that interaction with previously ejected material (either in the form of a stable wind or from erratic mass-loss episodes) is occurring at some level. The multiple outbursts of SN 2009ip detected in the 2009, 2011 and August 2012 (from 3 years to ∼ 1 month before the major 2012 explosion) are likely to have ejected conspicuous material in the immediate progenitor surroundings so that interaction of the 2012b explosion shock with this material qualifies as an efficient way to convert kinetic energy into radiation.
The radioactive decay of 56Ni represents another obvious source of energy. We employ the nebular phase formalism developed by Valenti et al. (2008) (expanding on the original work by Sutherland & Wheeler 1984 and Cappellaro et al. 1997) to constrain the amount of Nickel synthesized by the 2012b explosion using late time observations. If the observed light-curve were to be entirely powered by the energy deposition of the 56Ni radioactive decay chain, our latest photometry would imply MNi ∼ 0.03 M for a standard explosion kinetic energy of Ek = 1051 erg.48 For a low energy explosion with Ek = 1050 erg, MNi ∼ 0.08 M . Allowing for other possible sources of energy contributing to the observed luminosity (like interaction), we conclude MNi < 0.08 M . Using this value (and the photospheric formalism by Valenti et al. 2008, based on Arnett 1982) we largely underpredict the luminosity of SN 2009ip at peak for any value of mass and kinetic energy of the ejecta: the energy release of SN 2009ip is therefore not powered by 56Ni radioactive decay. Fraser et al. (2013) inde-
48 This limit is also sensitive to the ejecta mass Mej. We solve for the degeneracy between Mej and Ek using the observed photospheric velocity at maximum light vphot ∼ 4500 km s−1, which implies (Mej/M ) ∼ 3.0(Ek/1051 erg).

NIR emitting shell R"4x1015 cm

V!5000 km/s

Shock break out shell R!5x1014 cm

FIG. 26.— Sketch of SN 2009ip. Its environment as well as its ejecta are likely to have an extremely complex, potentially asymmetric structure. Here we show the basic components of the 2012b explosion and its environment. In the ejecta we recognize the presence of three velocity components at v ∼ 2500 km s−1, v ∼ 5000 km s−1 together with very fast material at v > 10000 km s−1 (Fig. 22). Shock break-out from a dense shell of material left over by the 2012a outburst is responsible for the major peak in the light-curve (Section 7.2). Material sitting at larger distance, connected with previous episodes of eruption, is responsible for partially re-processing the radiation into the NIR band, producing the NIR excess of Fig. 13 (Sec. 7.8). Our analysis requires this material to have a low filling factor and/or asymmetric distribution.
pendently derived MNi < 0.02 M , consistent with our findings. In the following we explore a model where the major UV-bright peak is powered by shock break-out from a dense shell ejected by the precursor bump, while continued interaction with previously ejected material is responsible for the peculiar, bumpy light-curve that follows.
7.2. Shock break-out plus continued interaction scenario for the 2012b explosion
The rapid rise and decay times of the major 2012b explosion (Fig. 25) suggest that the shock wave is interacting with a compact shell(s) of material. The relatively fast fading of CSM-like features and subsequent emergence of Type IIP features shown in Fig. 21 supports a similar conclusion. The bumps in the light-curve further suggest an inhomogeneous medium. We consider a model where the ejecta from the 2012b explosion initially interact with an optically thick shell of material, generating the UV-bright, major peak in the lightcurve (Fig. 26). In our model, the light-curve is powered at later times by interaction with optically thin material.
In the shock break-out scenario the escape of radiation is delayed with respect to the onset of the explosion until the shock is able to break-out from the shell at an optical depth τw ≈ c/vsh. This happens when the diffusion time td becomes comparable to the expansion time. Radiation is also released on the diffusion time scale, which implies that the observed bolometric light-curve rise time is trise ≈ td. Following Chevalier & Irwin (2011), the radiated energy at break-out Erad, the diffusion time td and the radius of the contact discontinuity at t = td (≡ Rbo) depend on the explosion energy E, the ejecta mass Mej, the environment density ρw (parametrized by the progenitor mass-loss rate) and opacity k. From our data we measure: trise ≈ 10 days; Rbo ≈ 5 × 1014 cm; Erad ≈ 1.3 × 1049 erg. We solve the system of equations for our observables in Appendix A and obtain the following estimates for the prop-

26 Margutti et al.

erties of the explosion and its local environment. Given the likely complexity of the SN 2009ip environment, those should be treated as order of magnitudes estimates.
The onset of the 2012b explosion is around 20 days before peak (2012 September 13). Using Eq. A2 and Eq. A3, the progenitor mass-loss rate is M˙ ≈ 0.07(vw/200 km s−1)M yr−1. We choose to renormalize the mass-loss rate to 200 km s−1, which is the FWHM of the narrow emission component in the Hα line (e.g. Fig. 16). The observed bolometric luminosity goes below the level of the luminosity expected from continued interaction of Eq. A6 around 32 days after the onset of the explosion or tpk + 12 days. By this time she shock must have reached the edge of the dense wind shell: tw 32 days. This constrains the wind shell radius to be Rw ≈ 1.2 × 1015cm (Eq. A7), therefore confirming the idea of a compact and dense shell of material, while the total mass in the wind shell is Mw ≈ 0.1M (Eq. A5).49 The system of equations is degenerate for Mej/E2. Adopting our estimates of the observables above and Eq. A1 we find Mej ≈ 50.5(E/1051erg)2M .50 The efficiency of conversion of kinetic energy into radiation depends on the ratio of the total ejecta to wind shell mass (e.g. van Marle et al. 2010; Ginzburg & Balberg 2012; Chatzopoulos et al. 2012). This suggests Mej ≈ Mw as order of magnitude estimate, from which E ∼ 1050 erg.
After tw the bolometric luminosity starts to decay faster, especially at UV wavelengths (Fig. 2). By this time the shock has overtaken the dense thick shell and starts to interact with less dense, optically thin layers of material producing continued power for SN 2009ip. In this regime the observed luminosity tracks the energy deposition rate: L = 4πR2ρw(vfs − vw)3, where R is the radius of the cold dense shell that forms as a result of the loss of radiative energy from the shocked region; vfs is the forward shock velocity; vw and ρw are the velocity and density of the material encountered by the shock wave (Chevalier & Irwin 2011 and references therein). The presence of clearly detected bumps in the bolometric light-curve (with associated rise in the effective temperature of the radiation, Fig. 11) suggests that the medium has a complex structure and it is likely inhomogeneous. consequently ρw might significantly depart from the ∝ R−2 profile expected in the case of steady wind. The increasing FWHM with time measured for the narrow component of the Hα line in Fig. 16 points to larger vw at larger distances from the explosion, therefore deviating from the picture of a steady wind with constant vw (see Section 7.3). Given the complexity of the explosion environment, we adopt a simplified shock interaction model (see e.g. Smith et al. 2010a) and parametrize the observed luminosity as: L = (η/2)wv3, where η is the efficiency of conversion of kinetic energy into radiation; w(R) ≡ M˙ /vw (hence ρw = w(R)/4πR2); while v is a measure of the expansion velocity of the shock into the environment. We estimate v from the evolution of the black-body radius with time (v ≈ R˙ HOT of Fig. 11), assuming that for tpk + 17 days the true shock radius continues to increase with v ≈ 4500 km s−1 (while the measured RHOT stalls and then decreases as the interaction shell progressively transitions to the optically thin regime, see Section 3). Using the bolometric
49 The mass swept up by the shock by the time of break-out is ≈ 0.05M . 50 Using the line of reasoning of Section 7.1, the relation between Mej and E just found implies MNi < 0.02 M for E = 1051 erg and MNi < 0.07 M for E = 1050 erg, consistent with the limits presented in Sec. 7.1.

luminosity of Fig. 11, we can therefore constrain the properties (density and mass) of the environment as sampled by the 2012b explosion.
We find that the total mass swept up by the shock from tw = 32 days until the end of our monitoring (112 days since explosion) is Mwthin ≈ (0.05/η)M . The total mass in the environment swept up by the 2012b explosion shock is therefore Mtot = Mw + Mwthin ≈ (0.2 − 0.3)M for η = 50 − 30%. As a comparison, Ofek et al. (2013a) derive a mass of ∼ 0.1M . Our analysis points to a steep density profile with ρw ∝ R−5.3 for R > 1.4 × 1015cm. The mass-loss rate is M˙ (R) = w(R)vw(R). We estimate vw from the evolution of the FWHM of the narrow Hα component in Fig. 16. Combining this information with the expression above we find M˙ (R) ∝ R−2 for (1.4 < R < 4.4) × 1015cm, with M˙ (R) ≈ (0.08/η)M /yr at R = 1.4 × 1015cm, declining to (0.008/η)M /yr at R = 4.4 × 1015cm.
7.3. Origin of the interacting material in the close environment
During the 2012b explosion, the shock interacts with an environment which has been previously shaped by the 2012a explosion and previous eruptions. In this section we infer the properties of the pre-2012a explosion environment, using the 2012a outburst as a probe. We look to: (i) understand the origin of the compact dense shell with which the 2012b shock interacted, whether it is newly ejected material by the 2012a outburst or material originating from previous eruptions; (ii) constrain the nature of the slowly moving material (v ≈ a few 100 km s−1) responsible for narrow line emission in our spectra.
We put an upper limit on the total amount of mass in the surroundings of SN 2009ip before the 2012a explosion assuming that the observed luminosity of the 2012a outburst is entirely powered51 by optically thin shock interaction with some previously ejected material of mass Mw12a. As before: L = (η/2)wv3. Mw12a = w(R)dR. The evolution of the blackbody radius of Fig. 11 suggests v ≈ 2500 km s−1 before 2012 August 21. We apply the same line of reasoning as above and assume that the shock continues to expand with this velocity, while the photosphere transitions to the thin regime and stalls at ≈ 0.4×1015 cm. In this picture, the 2012a shock sampled the environment on distances R < 1.2 × 1015cm, sweeping up a total mass of Mw12a ≤ (0.02/η)M . In Section 7.2 we estimated that the total mass of the dense wind shell from which the 2012b shock breaks out is Mw ≈ 0.1 M with radius Rw ≈ 1.2 × 1015cm. The wind shell mass is Mw = Mw12a(R < Rw) + Me1j2a(R < Rw), where Me1j2a(R < Rw) is the portion of the ejecta mass of the 2012a explosion within Rw at t = tw52. This implies Me1j2a(R < Rw) > 0.06 M (> 0.09 M ) for our fiducial efficiency η = 30% (η = 50%), comparable with the mass of the dense wind shell. We conclude that the UVbright 2012b explosion results from shock break-out from a dense shell which mostly (if not entirely) originates from the ejecta mass of the 2012a explosion, therefore establishing a
51 Any additional source of power would lower the required interacting mass.
52 Assuming MJD 56140 (2012 August 1) as zero-time for the 2012a outburst, this would correspond to ejecta with velocity ≤ 2000 km s−1 for free expansion.

SN2009ip

27

direct connection between the properties of the 2012a-2012b episodes.
The previous result also implies a solid lower limit on the total ejecta mass of the 2012a outburst: Me1j2a > 0.06 M (> 0.09 M ) for η = 30% (η = 50%). In Section 7.2 we estimated that the total mass collected by the 2012b shock by the end of our monitoring is Mtot ≈ 0.2 − 0.3 M , which constrains 0.06 < Me1j2a < 0.3 M for η ≥ 30%. In the following we use Me1j2a ≈ 0.1M as an order of magnitude estimate for the mass ejected by the 2012a outburst.53
Our spectra show evidence for narrow line emission (Section 4) typically observed in SNe IIn (and LBVs), which is usually interpreted as signature of the ejecta interaction with material deposited by the progenitor wind before explosion. For SN 2009ip we observe during the 2012b event a velocity gradient in the narrow emission from Hα (Fig. 16, panel c), with increasing velocity with time. This increase is consistent with being linear with time. This might suggest a Hubble-like expansion for the CSM following the simple velocity profile v ∝ R: as time goes by, the shock samples material at larger distances from the explosion (hence with larger velocity v). Our analysis indicates that episodes of mass ejection with approximate age 11-19 months before the 2012b explosion (roughly between February and October 2011) might reasonably account for the observed velocity gradient. We suggest that CSM material in the surroundings of SN 2009ip moving at velocities of hundreds km s−1 originates from this sudden episode(s) of mass ejection. Remarkably, SN 2009ip has been reported to be in eruptive phase between May and October 2011 (Pastorello et al. 2012), consistent with this picture.
The Hubble-like flow is not consistent with a steady wind and points instead to some mechanism leading to explosive mass ejections. Interestingly, it is during the September 2011 outburst that SN 2009ip showed evidence for material with unprecedented velocity, reaching v = 12500 km s−1 (Pastorello et al. 2012). Since no line-driven or continuum-driven wind mechanism is known to be able to accelerate stellar surface material to these velocities (Mauerhan et al. 2013), stellarcore related mechanisms have to be invoked. The explosive mass ejection suggested by our analysis might therefore be linked to instabilities developing deep inside the stellar core.
7.4. The role of asymmetries in SN 2009ip
The analysis of Section 4.5 indicates the presence of ejecta traveling at three distinct velocities: v ∼ −12000 km s−1, v ∼ −5500 km s−1 and v ∼ −2500 km s−1. These values correspond to the velocity of material seen in absorption (i.e. placed outside the photosphere). The radius of the hot blackbody RHOT of Fig. 11 tracks the position of the photosphere with time. Assuming free expansion of the ejecta and the explosion onset time (tpk −20 days) derived in the previous sections, we can predict at which time tv ejecta moving at a certain velocity v will overtake the photosphere at RHOT. Only for t tv can the ejecta give rise to absorption features in the spectra. Spherical symmetry is an implicit assumption in the calculation of RHOT, so that comparing the predicted tv to the observed time
53 Strictly speaking, we are only sensitive to the 2012a ejecta mass that has been overtaken by the 2012b explosion by the end of our monitoring. However, the analysis by Pastorello et al. (2012) shows evidence for strong deceleration of the 2012a ejecta by 2012 September 15, which suggests that most of Me1j2a has been encompassed by the 2012b explosion ∼ 100 days after.

FIG. 27.— Radio luminosity at peak vs. X-ray luminosity at the radio peak for a sample of Type IIn SNe (black dots). SN 2009ip is shown in red. Data have been collected from the literature. (See Pooley et al. 2002; Smith et al. 2007; Stritzinger et al. 2012; Chandra et al. 2012a; Pooley et al. 2007; Chandra & Soderberg 2007; Chandra et al. 2009; Zampieri et al. 2005; Houck et al. 1998; Chevalier 1987; van Dyk et al. 1993; Fabian & Terlevich 1996; and references therein).
of appearance of the absorption edges makes it possible to test the assumption of spherical symmetry of the explosion.
For v ∼ −2500 km s−1 we find tv = tpk + 55 days (2012 November 27) in excellent agreement with our observations, which constrain the v ∼ −2500 km s−1 absorption edge to appear between 2012 November 23 and December 5 (Fig: 22). No departure from spherical symmetry needs to be invoked for slow-moving ejecta, which likely includes most of the ejecta mass.54 Spherical symmetry is instead clearly broken by the high-velocity material traveling at v ∼ −12000 km s−1. For the 2012b explosion we detect high velocity material in absorption starting from ∼ 1 week after peak. Around peak the spectrum of SN 2009ip is optically thick and shows no evidence for material with v ∼ −12000 km s−1 (Fig. 19 and Fig. 20). However, in no way could a perfectly spherical photosphere traveling at 4000 − 5000 km s−1 mask the fast-moving ejecta at any time during the evolution, and in particular until the first week after peak, as we observed. This indicates a departure of the high-velocity ejecta from spherical geometry and might suggest the presence of a preferred direction in the explosion.
Asymmetry can also have a role in the spatial distribution of the interacting material, as supported by the observed coexistence of broad and (unresolved) intermediate components in the spectrum (Fig. 14). In this respect, Chugai & Danziger (1994) proposed the possibility of an enhanced mass loss on the equatorial plane of SNe IIn to explain the intermediate velocity component in SN 1988Z, other explanations being a clumpy circumstellar medium or, again, an asymmetric flow. In this context it is worth noting that asymmetry is also a likely explanation for the discrepant mass-loss estimates found by Ofek et al. (2013a), as noted by the authors. A disk-like geometry for SN 2009ip was proposed by Levesque et al. (2012) based on the H Balmer decrement. Finally, the binary-star merger scenario proposed by Soker & Kashi (2013) to interpret SN 2009ip naturally leads to ejecta with a bipolar structure.
54 Note however that in no way this argument can be used as a proof of spherical symmetry.

28 Margutti et al.

7.5. Hard X-rays and X-rays from shock break-out and
continued interaction
SN 2009ip is a weak X-ray and radio emitter (Fig. 27). In the following two sections we connect the lack of high X-ray/radio luminosity to the shock break-out plus interaction scenario we developed to explain the optical properties of SN 2009ip.
The shock break-out plus continued interaction scenario gives rise to a two-component spectrum, with a hard (X-rays) and a soft (UV/optical, at the break-out velocity of interest here) component (e.g. Svirski et al. 2012; Chevalier & Irwin 2012). The hard component is generated by bremsstrahlung emission from hot electrons behind the shock. Theory predicts that Lhard,bo ∼ 10−4Lbo (where Lbo is the break-out luminosity, resulting from the soft and hard component) as long as: (i) Inverse Compton (IC) cooling dominates over bremsstrahlung; (ii) high-energy photons undergo Compton degradation in the unshocked wind during their diffusion to the observer. We show in the following that both processes are relevant for SN 2009ip and provide a natural explanation for its very low X-ray to optical luminosity ratio (Lx/Loptical 10−4).
From our modeling of Section 7.2, we inferred a density parameter D∗ ≈ 0.4 (where D∗ is a measure of the density of the dense wind shell ρw, as explained in Appendix A). With a shock velocity v ≈ 4500 km s−1, the density measurement above implies that around break-out time the main source of cooling of the hot electrons is IC (see Chevalier & Irwin 2012). For our parameters, the IC to Bremsstrahlung (ff) emissivity ratio at break-out is (Svirski et al. 2012, their Eq. 17) IC/ f f ≈ 0.01(v/109cm s−1)2 or IC/ f f ≈ 0.05 for the observed v ≈ 4500 km s−1. IC is the dominant cooling source, suppressing the emission of hard photons in SN 2009ip. The calculations by Ofek et al. (2013a) instead assume negligible IC cooling.
Comptonization of the hard photons as they propagate through the unshocked wind region to the observer furthermore leads to a suppression of high-energy radiation. This process can effectively suppress photons with ≈keV energy if τes 15 − 20, the photon energy being limited by max = 511/τe2s keV. Our modeling of Section 7.2 implies τes ≈ 15 for SN 2009ip around shock break-out. This demonstrates that both the domination of IC over bremsstrahlung (i) and Compton losses (ii) are relevant to explain the weak X-ray emission in SN 2009ip. Identifying Lx with Lhard and Lbol with Lbo, the shock break-out scenario therefore naturally accounts for the observed Lx/Loptical 10−4 ratio even in the absence of photoabsorption.
Chevalier & Irwin (2012) calculate that full ionization (which gives minimal photo-absorption) is achieved for highvelocity shocks with v 104km s−1. For SN 2009ip v ≈ 4500 km s−1 likely leads to incomplete ionization (i.e. potentially important photo-absorption), that will further reduce the escaping X-ray flux. Using the explosion observables and Eq. A9 we constrain the total column density between the shock break-out radius and the observer to be N ≈ 2 × 1025cm−2 which gives a bound-free optical depth τbf ≈ 2 × 103 at 1 keV (Eq. A10).55 Since τbf ∝ R−1 and R ∝ t, soft (E ≈ 2 keV)
55 This calculation assumes a neutral medium and adopts the approximation by Ofek et al. (2013a) which is accurate within a factor of 2 with respect to more detailed calculations of the cross section by Morrison & McCammon (1983).

X-ray emission would not be expected from SN 2009ip until R ≈ 2 × 1016cm which happens for ∆t 440 days since the explosion, or 44 break-out time scales td (where td ≈ trise in our scenario).56 This calculation assumes an extended and spherically symmetric wind profile. The Chandra detection of SN 2009ip at much earlier epochs (∆t ≈ 4td) indicates that at least one of these assumptions is invalid, therefore pointing to a truncated and/or highly asymmetric wind profile. This is consistent with the picture of a dense but compact wind shell of radius Rw ≈ 1.2 × 1015cm followed by a steep density gradient ρw ∝ R−5.3 we developed in Section 7.2. Asymmetry also plays a role, as independently suggested by observations in the optical (Section 7.4).
Finally, Katz et al. (2011) predict that hard-X-ray emission with typical energy of 60 keV is also produced by the collision-less shock. The details of the spectrum are however unclear. Bound-free absorption is less important at these energies giving the chance to detect hard X-rays at earlier times. Our Swift-BAT campaign in the 15-150 keV range revealed a tentative detection. With these observations we put a solid upper limit on the hard X-ray to optical luminosity around maximum light which is LX,hard/Lbol < 5 × 10−3 at 5σ. The broad band SED around maximum light is shown in Fig. 28.
7.6. Radio emission from shock break-out and continued interaction
The shock-CSM interaction is a well known source of radio emission (e.g. Chevalier 1982; Chevalier 1984). The limited shock velocity we infer for SN 2009ip likely leads to a partially ionized medium as discussed above, so that freefree absorption plays a key role in suppressing the emitted radio flux. We quantify this statement below. The high density of the wind shell derived from our modeling of Section 7.2 implies a free-free optical depth at radius R, τff ≈ 3.8 × 1054(ν/GHz)−2.1(R/cm)−3 (Eq. A11). With τff ≈ 105 − 108 at R ≈ Rbo, no detectable radio emission is expected around break-out, consistent with our lack of radio detection around these times. If the dense wind profile extends out to large distances, the calculation above shows that no radio emission is expected until very late times, when the shock reaches R ≈ 7 × 1016cm. Our radio detection at much earlier epochs (∆t ≈ 60 − 80 days since explosion) demonstrates instead that the dense wind shell is not extended but truncated and adds further, independent evidence for a complex medium where inhomogeneity, asymmetry and/or low wind filling factor might also be relevant. This is consistent with the idea we suggested in Section 7.3 that the dense shell was the outcome of a short eruption, since eruptions are more likely to eject shells with sharp density edges as opposed to a steady mass loss.
7.7. GeV and neutrino emission at shock break-out
In the previous sections we suggested the presence of a dense (n ≈ 4 × 1010cm−3 at the break-out radius) and compact shell in the close environment of SN 2009ip that we associate with material ejected during the 2012a outburst. Shock break-out from this shell powers the major re-brightening of SN 2009ip at optical wavelengths (2012b explosion) and naturally explains the weak radio and X-ray emission we observe.
56 Adopting standard parameters Svirski et al. (2012) find that under standard parameters the X-ray emission is expected to dominate the energy release on time scales of the order of 10 − 50td.

SN2009ip

29

FIG. 28.— Broad-band SED of SN 2009ip around the optical peak. Shaded bands highlight different components of emission that dominate at different wavelengths. Neutral pion decay leads to γ-rays which are ultimately powered by the collision between the ejecta and extremely dense CSM shells of material.
We show here the limits obtained by Fermi-LAT in the week around maximum light (from t = tpk − 2 days until t = tpk + 4 days). Dashed purple line: expected GeV emission based on the shock explosion parameters derived in Section 7.2 and the model by Murase et al. (2011) assuming no attenuation. The thick-line model includes partial attenuation as relevant for a clumpy and/or asymmetric medium (Sec. 7.4) with filling factor < 1 that reduces the effects of the Bethe-Heitler
process. Soft and hard X-ray emission originates from the interaction of the shock with the CSM: both thermal and non-thermal emission might arise, mainly depending on the environment density. We show here a power-law (non-thermal emission, thick line with Fλ ∝ λ−2.4) and a thermal bremsstrahlung model with kT = 60 keV (dashed line) that fit the Swift-XRT 0.3-10 keV data obtained between tpk − 2 days and tpk + 4 days. Swift-BAT 15-150 keV data acquired at tpk ± 1 days are also shown. The dotted line marks the energy of 60 keV: this is the typical frequency of high-energy photons emitted by the collisionless shock that
forms at shock break-out (Katz et al. 2011). Optical and UV emission is powered by the shock break-out plus continued shock interaction with the environment. The orange thick line is the best-fitting "hot" black-body component with T ≈ 17000 K obtained at maximum light. NIR and IR emission could originate from dust formation, dust vaporization but it might also be a light echo. Due to the high temperature we infer (T ≈ 3000 K, red solid line) we favor dust vaporization
here. Shock-CSM interaction is also the source of millimeter and radio photons through synchrotron emission. At these times this emission was quenched by
free-free absorption by the dense shell of material.

FIG. 29.— Predicted muon and anti-muon neutrino fluence from SN 2009ip
using the observables and explosion parameters of Sec. 7.2 according to the
model by Murase et al. (2011). The shaded band corresponds to different val-
ues for the total energy in cosmic rays ECR here parametrized as a function of the radiated energy at trise. Note that, given Ek Erad, larger values of iEnCteRgraarteedalosovepro∆sstib≈let.risDe a(szheendithb-launegllienea:veartamgeodspwheitrhicinn1e◦u)t,riwnohibchaciksgsrhoouwnnd for comparison. For this event, the atmospheric neutrino background is more
severe since SN 2009ip occurred in the southern hemisphere. For better lo-
calized explosions, this plot shows how limits on the neutrino emission can
be used to constrain ECR.

The collision of the ejecta with massive shells is also expected to accelerate cosmic rays (CRs) and generate GeV gammarays (Murase et al. 2011; Katz et al. 2011) with fluence that depends both on the explosion and on the environment parameters. The GeV emission is expected to be almost simultaneous with the UV-optical-NIR emission. The proximity of SN 2009ip (≈ 24 Mpc) justifies the first search for GeV emission from shock break-out. Fermi-LAT observations covering the period 2012 September 3 – October 31 (tpk − 30 days until tpk + 28 days) yielded no detection (Sec. 2.10). Following Murase et al. (2011), we predict the GeV emission from SN 2009ip (2012b outburst) using the explosion and environment parameters inferred in Sec. 7.2 (shell density ρw or mass Mw) together with the observables of the system (trise, Rbo). We take into account the attenuation due to γγ → e+e− pair production, assuming a black-body spectrum with T = 20000 K. The attenuation by extragalactic background light is also included. Since n is not too large, the Bethe-Heitler pair production will be irrelevant in our case, as it happens for a clumpy or asymmetric CSM (Sec. 7.4) with filling factor < 1. The injected CR energy ECR is assumed to be equal to that of the radiation energy at trise. We compare the expected GeV fluence with observations in Fig. 28: the Fermi-LAT non-detection is consistent with our picture of ejecta crashing into a compact and dense but low-mass ( 0.1 − 0.2 M ) shell of material at small radius that we developed in the previous sections. For the detection of γ−rays, brighter SNe (closer SNe or SNe accompanying larger dissipation) are needed.
Using the set of parameters above we predict the muon and

30 Margutti et al.

anti-muon neutrino fluence from SN 2009ip in Fig. 29, using the model by Murase et al. (2011). CRs produce mesons through inelastic proton-proton scattering, leading to neutrinos as well as γ−rays. As for the γ−rays, given that the explosion/ environment parameters are constrained by observations at UV/optical/NIR wavelengths, the neutrino fluence directly depends on the total energy in CRs, ECR (or, equivalently on the CR acceleration efficiency CR, ECR ≡ CRE, being E the explosion energy). Note that the injected CR energy ECR is assumed to be the radiation energy at trise, but larger values of ECR are also possible. The maximum proton energy is also set to 1.5 PeV, corresponding to εB = 0.01. Our calculations show that SN 2009ip was a bright source of neutrinos if ECR is comparable to the radiated energy. We note however that its location in the souther sky was not optimal for searches for neutrino emission by IceCube because of the severe atmospheric muon background. With this study we demonstrate how it is possible to constrain the CR acceleration efficiency if: (i) the explosion/ environment parameters are constrained by observations at UV/optical/NIR wavelengths; (ii) deep limits on muon neutrinos are available, as will be the case also for southern sky sources once KM3Net will be online in the near future.
7.8. Origin of the NIR excess
The time-resolved broad-band SED analysis of Section 3 identifies the presence of a NIR excess of emission (Fig. 13) and allows us to constrain its temporal evolution (Fig. 11). Contemporaneous NIR spectroscopy from Section 2.6 clearly shows that the NIR excess cannot be ascribed to line emission (Fig. 7), therefore pointing to a physical process producing NIR continuum emission. Adopting a black-body spectral model (with the implicit assumption of spherical symmetry) we find that: (i) the equivalent black-body radius is RCOLD ∼ 4 × 1015 cm with very limited evolution with time over 30 days of monitoring.57 This is in contrast with the hot component radius RHOT which increases linearly with time inside the cooler component. (ii) The cold black-body temperature is also stable, with TCOLD ∼ 3000 K (while the hot black body cools from 19000 K to < 10000 K). Considering that our fits can overestimate the true dust temperature of hundreds of degrees (up to 20% according to Nozawa et al. 2008) the real dust temperature might be close to ∼ 2500 K. (iii) The resulting NIR excess luminosity is LCOLD ∼ 4 × 1041erg s−1 which represents (2 − 4)% of LHOT. We use these properties to constrain the origin of the NIR excess below.
A clear spectroscopic signature of dust formation is the development of highly asymmetric and blue-shifted line profiles (see e.g. Smith et al. 2008, their Fig. 4) which is not observed. The temperature of ∼ 2500 − 3000 K is also prohibitively high for dust to form. At this very early epochs the shock radius is also < RCOLD. Note that RCOLD, derived assuming a blackbody spectrum, represents a lower limit to the real size RNIR of the NIR emitting region (Fig. 26): RNIR ∝ RCOLD/ f 0.5, where f < 1 is the covering factor of the NIR emitting material. f < 1 is indeed required for the hot black-body radiation to be able to escape. This clearly implies that the NIR emission cannot originate from dust created behind the reverse shock. The dust creation scenario is therefore highly unlikely, leading to the conclusion that the NIR excess originates from preexisting material ejected by the progenitor before the 2012b
57 Our observations imply vCOLD < 103 km s−1.

explosion. Since the geometry of the NIR emitting material can be non spherical (as we find in Section 7.4) and/or have low filling factor (i.e. the material can be clumpy), we do not expect this material to necessarily produce absorption along our line of sight.
Material ejected during the 2012a outburst would be required to travel at an average velocity v > 8000 km s−1 (to reach 4 × 1015cm at the time of the 2012b explosion) and then to decelerate to v < 1000 km s−1 (to match the observed evolution of RCOLD). We consider this scenario unlikely58 and suggest that the origin of the pre-existing material is rather linked to the eruption episodes of the progenitor of SN 2009ip in the years before. The same conclusion was independently reached by Smith et al. (2013). We note that the size of the cool emitting region of SN 2009ip is remarkably similar to the pre-outburst dust radii of other optical transients linked to eruption episodes of their progenitor stars like NGC 300 OT2008-1 (R ∼ 5 × 1015 cm) and SN 2008S (R ∼ 2 × 1015 cm, see e.g. Berger et al. 2009, their Fig. 28).
NIR emission in Type II SNe has also been connected to the extended atmosphere of the expanding star (e.g. Dwek et al. 1983). However, for SN 2009ip, the large radius we infer from our modeling (RCOLD ∼ 4 × 1015 cm) and the lack of a clear evolution of the temperature with time argue against this interpretation.
A light echo from dust (i.e. pre-existing dust heated up by the UV and optical radiation from the explosion) would require the dust grains to survive the harsh environment. At the high temperature of T ∼ 2500 − 3000 K this is however unlikely, while dust vaporization is more likely to happen (see e.g. Draine & Salpeter 1979). We speculate on the dust vaporization scenario below.
In this picture a cavity is excavated by the explosion radiation out to a radius Rc: this radius identifies the position of the vaporized dust, while it does not track the outer dust shell radius (which is instead likely to expand with time, see e.g. Pearce et al. 1992). Being the dust shell created in the years before the 2012 explosion, we expect the smaller grains to be located at the outer edge of the dust shell as a result of forces acting on them (e.g. radiative pressure). At smaller distances we are likely dominated by the larger dust grains. Following Dwek et al. (1983), their Eq. 8 (see also Draine & Salpeter 1979) the radius of the dustfree cavity for a UV-optical source with luminosity L is: Rc = 23( Q (L/L )/(λ0/µm)T5)0.5, where Rc is in units of pc; T is the temperature at the inner boundary of the dust shell which we identify with the evaporation temperature Tev; Q is the grain emissivity: Q = (λ0/λ)n for λ ≥ λ0 while Q = 1 for λ < λ0. Here we adopt n = 1.59 The dust grain radius is a, with λ0 ∼ 2a which implies Rc ∝ a−0.5. For Rc = RCOLD ∼ 4 × 1015cm, Tev ∼ 3000 K and L = Lpk ∼ 1.7 × 1043erg s−1, we constrain the radius of the vaporizing dust grains to be a (2 − 5)µm.
Grain radii of 0.2 − 2 µm are typically found in dust shells, suggesting that we are possibly witnessing the vaporization of the larger grains at R ∼ Rc. The very high vaporization temperature is only potentially compatible with materials like
58 Our SED fitting indicates the presence of a NIR excess with similar radius during the 2012a eruption as well (Fig. 11), adding further evidence that the NIR emitting material pre-existed the 2012a episode.
59 A value of n between 1 and 2 is usually assumed. Using n = 2 have no impact on our major conclusions. In particular, our estimate of the dust grain radius would be a ≈ 1µm.

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medium, previously shaped by multiple episodes of shell ejection by the progenitor at different times. Here we address three major questions:
• What is the nature of the SN 2009ip double explosion in 2012?
• What is the underlying physical mechanism?
• What is the progenitor system of SN 2009ip?
We address the first two questions by considering the close similarity of SN 2009ip and SN 2010mc (Section 8.1), the properties of the progenitor system of SN 2009ip (Section 8.2), the physical mechanisms that can lead to sustained mass loss (Section 8.3) and the constraints on the energetics we derived from our modeling (Section 8.4).

FIG. 30.— The comparison of the absolute R band magnitude of SN 2009ip and Type IIn SN 2010mc (Ofek et al. 2013b) reveals a striking similarity between the two explosions both during the precursor-bump and the major outburst.
graphite, silicates, corundum and carbide, that have binding energy U0 6 eV. Even in these cases Tev ∼ 2500 − 3000 K requires extremely short vaporization time scales of the order of ∼ 1 day or less, and large grain dimensions of the order of 1 µm (see Draine & Salpeter 1979, their Eq. 24).
The passage of a SN explosion shock through a dust shell is one of the processes believed to establish the radius distribution of dust particles in our Universe. The explosion shock will eventually interact with the NIR emitting "shell" likely causing a flattening in the light-curve decay. This flattening should be then followed by a rapid decline once the shock reaches the edges of the NIR "shell". The timing of the interaction is however critically dependent on the velocity of the shock, the asymmetry of the explosion, the dust shell filling factor and/or the asymmetry of the dusty material (but also on the expansion velocity of the shell).60 Allowing for a deceleration of the blast wave to v ∼ 2500 km s−1 and a low preexisting dust filling factor f = 0.3, we anticipate the shock dust shell interaction to happen in the second half of 2013. SN 2009ip might be one of the rare cases where it has been possible to map the properties of dust before and after the interaction with the explosion shock. Future IR observations of SN 2009ip are of primary importance in this respect and will clarify how newly condensed dust in the explosion ejecta mixes with pre-existing dust.
Finally, we end by noting that a NIR echo from pre-existing gas can possibly account for the high temperature of the NIR excess while naturally explaining the almost flat JHK lightcurve around maximum light , compared to the steeply decaying UV emission of Fig. 2. This possibility is further explored in a dedicated paper (Margutti et al. in preparation). A complementary discussion of the NIR properties of SN 2009ip can be found in Smith et al. (2013) who favor the infrared echo hypothesis.
8. DISCUSSION
Our analysis characterizes the 2012b episode as a lowenergy, asymmetric explosion happening in a complex
60 High-velocity ejecta with v ∼ 10000 km s−1 reached RCOLD at tpk + 26 days, which is interestingly close to the time of the second major peak in the bolometric light-curve in Fig. 11.

8.1. SN 2009ip and SN 2010mc
In interpreting the fate and nature of SN 2009ip we have to consider its close likeness to SN 2010mc (Ofek et al. 2013b), the only other hydrogen-rich explosion with clear signs of interaction and a detected precursor before the major event. The similarity extends to the energetics, time scales (Fig. 30) and spectral properties both during the precursor bump and the major re-brightening, as noted by Smith et al. (2013).
The first important conclusion is that (i) the precursor and the major re-brightening are causally connected events, being otherwise difficult to explain the strictly similar phenomenology observed in SN 2009ip and SN 2010mc, two distinct and unrelated explosions. The same conclusion is independently supported by the very short time interval between the precursor and main outburst when compared to the progenitor star lifetime, as pointed out by Ofek et al. (2013b) for SN 2010mc. A second conclusion is that (ii) whatever causes the precursor plus major outburst phenomenology, this is not unique to SN 2009ip and might represent an important evolutionary channel for massive stars.
Furthermore, (iii) SN 2009ip and SN 2010mc must share some fundamental properties. In particular their evolution through the explosive phase must be driven by few physical parameters. A more complicated scenario would require unrealistic fine tuning to reproduce the close similarity of SN 2009ip and SN 2010mc. This also suggests we are sampling some fundamental step in the stellar evolution of the progenitor system. Finally, (iv) whatever the physical mechanism behind, the time interval of ∼ 40 days (Fig. 30) between the precursor explosion and the main event must be connected to some physically important time scale for the system.
We employ the fast χ2 technique for irregularly sampled data by Palmer (2009) to search for periodicity and/or dominant time-scales in the outburst history of SN 2009ip before61 the major 2012 explosion. Details can be found in Appendix B. Applying the method above to the R-band data we find evidence for a dominant time-scale of ∼ 38 days, (with significant power distributed on time-scales between 30 and 50 days), intriguingly similar to the ∆t ∼ 40 days between the precursor and major explosion in 2012. We emphasize that this is not a claim for periodicity, but the identification of a dominant variability time-scale of the signal.
Our analysis identifies the presence of a fundamental timescale which regulates both the progenitor outburst history and
61 See Martin et al., in prep. for a temporal analysis of SN 2009ip during the main episode of emission in 2012.

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the major explosion, and that is also shared by completely independent events like SN 2010mc. This time-scale corresponds to a tiny fraction of a massive star lifetime: ∼ 10−8 for τ ∼ 4 − 6 Myr, as appropriate for a 45 − 85 M star. We speculate on the nature of the underlying physical process in the next two sections.
8.2. The progenitor system of SN 2009ip
In Section 7.4 we showed that asymmetry plays a role in the 2012 explosion, which might point to the presence of a preferred direction in the progenitor system of SN 2009ip. This suggests either a (rapidly?) rotating single star or an interacting binary as progenitor for SN 2009ip. We first update previous estimates of the progenitor mass of SN 2009ip using the latest stellar evolutionary tracks and then discuss the effects of stellar rotation and the possibility of a binary progenitor.
From HST pre-explosion images, employing the latest Geneva stellar evolutionary tracks (Ekström et al. 2012) which include important updates on the initial abundances, reaction rates and mass loss prescriptions with respect to Schaller et al. (1992), we determine for MV = −10.3 ± 0.3 (Foley et al. 2011) a ZAMS mass M 60 M assuming a non-rotating progenitor at solar composition62, consistent with previous estimates. MV = −10.3 corresponds to LV ∼ 106.1L . This implies Lbol > 2 × 106 L , thus rivaling the most luminous stars ever discovered (e.g. Crowther et al. 2010). Luminosities of a few 106 L have been indeed associated to the group of LBV stars with typical temperature of ∼ 15000 − 25000 K. Adopting this range of temperature results in Lbol ∼ 5 × 106 L , suggesting that any progenitor with M < 160 M was super Eddington at the time of the HST observations. For (2 < Lbol < 5) × 106 L the allowed mass range is 60 M < MZAMS < 300 M (e.g. Crowther et al. 2010). Including the effects of axial rotation results instead in a more constrained range of allowed progenitor mass: 45 M < MZAMS < 85 M (for Ω/Ωcrit = 0.4).
Rapid rotation strongly affects the evolution of massive stars, in particular by increasing the global mass-loss rate (e.g. Maeder & Meynet 2000). More importantly, Maeder (2002) showed that the mass loss in rapidly rotating massive stars does not remain isotropic, but it is instead enhanced at the polar regions, thus favoring bipolar stellar winds (e.g. Georgy et al. 2011, their Fig. 2). As a result, the formation of an asymmetric (peanut shaped) nebula around rapidly rotating stars is very likely. Additionally, rapid rotation can induce mechanical mass loss, resulting in some matter to be launched into an equatorial Keplerian disk. It is clear that any explosion/eruption of the central star will thus naturally occur in a non-isotropic medium, as we find for SN 2009ip. Rotation further leads to enhanced chemical mixing (e.g. Chatzopoulos et al. 2012; Yoon et al. 2012).
HST pre-explosion images cannot, however, exclude the presence of a compact companion.63 The massive progenitor of SN 2009ip might be part of a binary system and the repetitive episodes of eruption might be linked to the presence of an interacting companion. The close periastron passages of a companion star in an eccentric binary system has been invoked by Smith & Frew (2011) to explain the brightening episodes of ηCar in 1838 and 1843. In the context of
62 The simulations extend up to ZAMS mass of 120 M . 63 The variability argument allows us to conclude that the observed emission is however dominated by the progenitor of SN 2009ip.

SN 2009ip, Levesque et al. (2012) discussed the presence of a binary companion, while Soker & Kashi (2013) suggested that the 2012b explosion was the result of a merger of two stars, an extreme case of binary interaction. A binary scenario was also proposed by Soker (2013) for the cosmic twin of SN 2009ip, SN 2010mc.
A binary system would have a natural asymmetry (i.e., the preferred direction defined by the orbital plane) and a natural time scale (i.e., the orbital period) as indicated by the observations. A possible configuration suggested to lead to substantial mass loss is that of a binary system made of a compact object (NS) closely orbiting around a massive, extended star (Chevalier 2012 and references therein). In this picture the mass loss is driven by the inspiral of the compact object in the common envelope (CE) evolution and the expansion velocity of the material is expected to be comparable to the escape velocity for the massive star. For64 v ∼ 1000 km s−1 and M ∼ 45 − 85 M the radius of the extended star is R∗ < (1 − 2) × 1012cm (where the inequality accounts for the fact that we used the ZAMS mass). In this scenario, the mass loss is concentrated on the orbital plane of the binary (Ricker & Taam 2012), offering a natural explanation for the observed asymmetry (Section 7.4). However, it remains unexplained how this mechanism would be able to launch material with v ∼ 104 km s−1 as observed during the 2012a episode (Mauerhan et al. 2013, Pastorello et al. 2012). A potential solution might be the presence of a high-velocity accretion disk wind.While the presence of a dominant time scale common to the eruptive and explosive phases makes the binary progenitor explanation particularly appealing65, the common envelope model in its present formulation seems to have difficulties in explaining the observed phenomenology. Alternative scenarios are explored in Section 8.3.
8.3. Physical mechanisms leading to substantial mass loss in evolved single stars
The outer envelope of an evolved massive star contains ∼ 95% of the stellar radius but only a tiny fraction of the stellar mass. The physical mechanism leading to the 2012a outburst must have been fairly deep seated to unbind more than ∼ 0.1 M . It is furthermore required to generate a large amount of energy (1048 erg plus the energy to lift the mass out of the deep gravitational potential well) and to explain the presence of a dominant time scale of ∼ 40 days which is also shared with the observed eruption history in the years before the explosion. The mass-loss rate of at least a few ∼ 0.1 M /yr can only be sustained for an extremely small fraction of the life of a star and in principle only for ∼ years, before inducing important adjustments to the stellar structure. We explore here different physical mechanisms that have been proposed to lead to substantial mass loss in the late stages of evolution of massive single stars and discuss their relevance for SN 2009ip. They can basically be grouped into 4 categories: (i) pulsational pair-instability; (ii) super-Eddington fusion luminosity; (iii) Eddington-limit instabilities; (iv) shock heating of the stellar envelope.
Pulsational pair-instability (PPI, Barkat et al. 1967) applies to massive stars with helium core of ∼ 40 − 60 M . PPI leads to partial unbinding of the star, with a sequence of eruptions
64 This is the velocity required to reach Rbo after ∼ 50 days since the 2012a explosion.
65 Adopting t ∼ 40 days as orbital period implies an orbital radius R < 1013 cm.

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accompanied by the ejection of shells of material of the stellar envelope. This mechanism has been considered as a plausible explanation for SN 2009ip by Pastorello et al. (2012), Mauerhan et al. (2013) and Fraser et al. (2013) but has been rejected by Ofek et al. (2013b) and Ofek et al. (2013a) for both SN 2010mc and SN 2009ip (the leading argument being that PPI would lead to the ejection of much larger amounts of mass than the ∼ 10−1 − 10−2 M they infer). According to Woosley et al. (2007), a non-rotating star with zero metallicity (Z=0) and ZAMS mass 95 and 130 M meets the criteria for PPI. However, with updated prescriptions for the mass-loss rate, and assumption of solar metallicity (which more closely represents our conditions of 0.4 Z < Z < 0.9 Z , Section 5) Ekström et al. (2012) predict that even stars with M = 120 M will end the C-burning phase with mass ∼ 31 M , below the threshold of ∼ 40 M to trigger PPI. While our limits on the progenitor mass of SN 2009ip in Sec. 8.2 do not formally rule out the PPI scenario in the non-rotating case, they definitely allow for progenitors starting with much lower mass than required for the PPI to develop (M 120 M ).
Rapidly rotating progenitors (Ω/Ωcrit > 0.5) enter the PPI regime starting with substantially lower ZAMS mass: ∼ 50 M for Z = 0 (Chatzopoulos & Wheeler 2012; Yoon et al. 2012). On the other hand, for Z = Z even very massive rotating 120 M stars develop a core with M ∼ 19 M , insufficient to trigger PPI (Ekström et al. 2012). In the previous sections we constrained the progenitor of SN 2009ip with ZAMS mass 45 M < M < 85 M and metallicity 0.4 Z < Z < 0.9 Z . Following the prescriptions from Chatzopoulos & Wheeler (2012), adopting the most favorable conditions (i.e. M = 85 M and Z = 0.4 Z ) and starting with a very rapidly rotating star with Ω/Ωcrit = 0.9, we find the final oxygen core mass to have M ∼ 35 M , formally below the threshold of 40 M for PPI. This indicates some difficulties for rotating progenitors to reach the PPI threshold. What remains difficult to interpret both for rotating and non-rotating progenitors is the presence of a dominant time-scale of ∼ 40 days: depending on the pulse properties and the physical conditions of the surviving star, the interval between pulses can be anywhere between days to decades (Woosley et al. 2007).
Alternatively, convective motions stimulated by the superEddington fusion luminosity in the core of a massive star can excite internal gravity waves (g-mode non-radial oscillations, Quataert & Shiode 2012). An important fraction of the energy in gravity waves can be converted into sound waves, with the potential to unbind several M . For a 40 M star, Quataert & Shiode (2012) estimate that 1047 − 1048 erg will be deposited into the stellar envelope. If this mechanism is responsible for the ejection of material by the 2012a explosion, the velocity v ∼ 1000 km s−1 of the unbound material inferred from our observations implies an ejecta mass 0.1 M , consistent with our estimate of the mass of the compact shell encountered by the 2012b explosion in Section 7.2. However, this mechanism is likely to lead to a steady mass-loss (E. Quataert private communication) as opposed to shell ejection and does not offer a natural explanation for the 40-day time scale involved in the process. We therefore consider the super-Eddington luminosity mechanism unlikely to apply to SN 2009ip and SN 2010mc, (see however Ofek et al. 2013b).
With R-band observed magnitudes between ∼ 18 mag and ∼ 21 mag reported by Pastorello et al. (2012), SN 2009ip oscillates between ∼Eddington and super-Eddington luminosity episodes in the years preceding the double explosion in

2012. Depending on the effective temperature of the emission, R ∼ 21 mag corresponds to the Eddington luminosity of a star with mass 40 − 80 M . R ∼ 18 would correspond to an Eddington limit for a star of M = 600 − 1200 M . A number of instabilities have been shown to develop in stars approaching and/or exceeding the Eddington limit, actually allowing the super-Eddington luminosity to persist (e.g. Owocki & Shaviv 2012), potentially powering LBV-like eruptions. The eruption time scale, repetition rate and ejected mass are however loosely predicted, so that it is unclear if any of these mechanisms would apply to the eruption history of SN 2009ip and, even more importantly, how these are connected with the 2012 double explosion (which clearly differs from superEddington powered winds). With luminosity between ∼ 106 and ∼ 107 L the progenitor of SN 2009ip falls into the region where radiation pressure starts to have a major role in supporting the star against gravity (see e.g. Owocki & Shaviv 2012, their Fig. 12.1): we speculate that the dominance of radiation pressure over gas pressure in the envelope might have an important role in determining the repeated ejection of shells of SN 2009ip.
Indeed, the outer layers of massive stars close to the Eddington limit and/or critical rotation are only loosely bound and might be easily ejected if enough energy is suddenly deposited. A potential source of energy deposition has been identified in thermonuclear flashes associated with shell burning (Dessart et al. 2010). Differently from the PPI, this scenario does not require the progenitor to be extremely massive. Alternatively, non-radial gravity mode oscillations above the core or near to the H-burning shell could provide fresh fuel, triggering a burst of energy (and subsequent shell ejection) in massive stars like ηCar (Guzik 2005). Both scenarios have the advantage to be basically driven by two parameters: the deposited energy and the envelope binding energy, naturally satisfying the "simplicity" criterion of Section 8.1.
Finally we mention that direct numerical simulations of pre-collapse hydrodynamics by Meakin (2006), Meakin & Arnett (2007) and Arnett & Meakin (2011a,b) found eruptive instabilities due to the interaction of oxygen and silicon burning shells. The instabilities are related to turbulent convection (and being inherently nonlinear, are invisible to conventional linear stability analysis). These simulations specifically predicted mass ejection prior to core collapse, suggesting that pre-collapse evolution is far more eventful than previously thought. However, as for all the other mechanisms analyzed in this section, it is at the moment unclear how to explain the 40-day time scale.
8.4. The nature of SN 2009ip double-explosion in 2012: the explosive ejection of the stellar envelope
Our modeling shows that the 2012b explosion is not powered by Ni radioactive decay (Section 7.1) and demonstrates that a shock breaking out from a shell of material previously ejected by the 2012a outburst can reasonably account for the observed properties of the 2012b explosion, constraining the mass of the ejecta to Mej = 50.5(E/1051erg)2 M (Section 7.2). This strongly suggests an explosion energy well below 1051erg (likely around 1050 erg), and brings to question the fate of SN 2009ip: a terminal SN explosion (Mauerhan et al. 2013) or a SN "impostor" (Pastorello et al. 2012; Fraser et al. 2013)?
A key prediction of the SN scenario is the presence of chemical elements produced by the SN nucleosynthesis (like

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FIG. 31.— Our spectrum taken in April 2013 (tpk + 190 days) shows minimal evolution with respect to ∼ 100 days before. The most notable difference with respect to spectra acquired after or during the LBV-like eruptions is the presence of intermediate He I and [CaII] in emission. We still do not find evidence for
SN-synthesized material (e.g. [OI] 6300 Å ). We refer to Sahu et al. (2006) for details on Type IIP SN 2004et.

oxygen) in late-time spectra, which is not expected in the case of a non-terminal explosion. Our latest spectrum acquired on 2013 April 11 (tpk + 190 days) still shows no evidence for SNsynthesized material (Fig. 31). This non-detection is consistent with (but surely not a proof of) the non-terminal explosion scenario. This spectrum, dominated by H emission lines and where the high-velocity absorption features have completely disappeared, shares some similarities with preexplosion spectra of SN 2009ip. A notable difference, however, is the presence of (intermediate width) He I and [Ca II]. The present data set therefore does not yet offer compelling evidence for a standard (i.e. resulting from the collapse of a degenerate core) SN-explosion scenario.
Interaction seems to be still dominating the emission at this time. Our late-time observations indicate a flattening in the light-curve (our latest V-band photometry acquired at tpk + 190 days implies a decay of ∆V = +0.86 ± 0.22 over 98.5 days, or 0.0087 ± 0.0022 mag/day), which might be caused by the interaction of the shock with the NIR emitting material. Our model predicts that the “flat” phase should be followed by a rapid decline, once the shock has reached the edge of the “NIR shell”.
We propose that SN 2009ip was the consequence of the explosive ejection of the envelope of a massive star, which later interacted with shells of material ejected during previous eruption episodes. While E < 1051erg is insufficient to fully unbind a massive star, its outer envelope is only loosely bound, and can be easily ejected by a lower-energy explosion. The origin of the energy deposition is not constrained as unclear is its potential relation with the possible binary nature of the system. The final fate of SN 2009ip depends on the properties of the remaining “core”, and in particular, on its mass and rotation rate: if the star managed to explode its entire H envelope still retaining a super-critical core mass, it

might re-explode in the future as a genuine H-poor (Ib/c) SN explosion; if instead the star partially retained its envelope, it will possibly give rise to other SN-impostor displays, on time scales which are difficult to predict.66. The core might also have directly collapsed to a black hole: in this case this would mark the “end”. Only close inspection of the explosion site after SN 2009ip has appreciably faded will reveal if the star survived the ejection of its outer layers. Given the impressive similarity with SN 2010mc, our interpretation extends to this event as well.
9. SUMMARY AND CONCLUSIONS
The “2012 double-explosion” of SN 2009ip brought to light the limits of our current understanding of massive star evolution, pointing to the existence of new channels for sustained mass-loss in evolved stars, whose origins has still to be identified. Our extensive follow-up campaign and modeling allow us to identify the properties of a complex ejecta structure and an explosion environment shaped by repeating outbursts during the previous years. We find that:
• SN 2009ip is embedded in a sub-solar metallicity environment (0.4 Z < Z < 0.9 Z ) in the outskirts of its host galaxy (d ∼ 5 kpc), radiated ∼ 3 × 1049 erg during the 2012b explosion, reaching a peak luminosity of ∼ 2 × 1043 erg s−1. The 2012a precursor bump released ∼ 2 × 1048 erg.
• The explosion is not powered by 56Ni radioactive decay. From late-time photometry, the total 56Ni mass is MNi < 0.08 M . Narrow emission lines in the optical spectra require interaction of the explosion shock
66 Note that SN 2009ip might still undergo PPI in the future even if it managed to eject its entire envelope during the 2012b explosion.

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with the CSM at some level. We suggest this interaction is also responsible for mediating the conversion of the shock kinetic energy into radiation, powering the observed light-curve.
• Spectroscopy at optical to NIR wavelengths further identifies three distinct velocity components in the ejecta with v ∼ 12000 km s−1, v ∼ 5500 km s−1 and v ∼ 2500 km s−1, arguing against a continuous velocity distribution. No departure from spherical symmetry needs to be invoked for the slow-moving (v ≤ 5500 km s−1) ejecta. Instead, spherical symmetry is clearly broken by the high-velocity material, possibly pointing to the presence of a preferred direction in the explosion. Broad and intermediate components in the optical/NIR spectra are consistent with the view that asymmetry might also have a role in the spatial distribution of the interacting material.
• CSM material in the region surrounding SN 2009ip, with a velocity of a few to several 100 km s−1, originates from abrupt episode(s) of mass ejection in the previous years. Mass loss is unlikely to have occurred in the form of a steady wind: instead, our analysis favors explosive mass ejections, likely linked to instabilities developing deep inside the stellar core.
• We interpret the major 2012 re-brightening to be caused by the explosion shock breaking out through a compact and dense CSM shell previously ejected in the 2012a outburst. Our analysis constrains the onset of the explosion to be ∼ 20 days before peak. The breakout radius is Rbo ∼ 5 × 1014cm. The shell extends to Rw ∼ 1.2 × 1015cm, with a total mass of Mw ∼ 0.1 M . The presence of a compact and dense shell is independently supported by (i) the detection of X-ray radiation with Lx/Lopt < 10−4 and (ii) by the late-time rise of radio emission. After break-out the optical-UV lightcurve is powered by continued interaction with optically thin material characterized by a steep density profile.
• A shock breaking out through a dense shell is expected to produce hard X-rays with a typical energy of 60 keV. We report a tentative detection of hard X-rays (15-150 keV range) around the optical peak and put a solid upper limit on the hard X-ray to optical emission from SN 2009ip: LX,hard/Lbol < 5 × 10−3.
• The collision of the ejecta with a massive shell(s) is expected to accelerate cosmic rays and generate GeV gamma-rays. Using Fermi-LAT data we detect no GeV radiation, consistent with the picture of ejecta crashing into a compact and dense but low-mass shell of material at small radius.
• Our calculations indicate that the latest outburst from SN 2009ip was a bright source of neutrinos if the CR energy is comparable to the radiated energy. We demonstrate with SN 2009ip how to constrain the cosmic-ray acceleration efficiency using broad band electromagnetic data and deep limits on muon neutrinos that will be available in the near future when new facilities will come online.

• We speculate that the NIR excess of emission is a signature of vaporization of dust grains with radius (2 − 5)µm in a shell of material located at R > 4 × 1015cm. The very high temperature associated with the NIR excess remains however difficult to explain. The shell was ejected by the progenitor in the years preceding 2012, and will be overtaken by the explosion shock during 2013.
• Finally, our modeling of the 2012b explosion implies a small ejecta mass Mej ∼ 0.5 M , with an explosion energy well below 1051 erg (likely around 1050 erg), thus raising questions about the nature of SN 2009ip: was it a terminal SN explosion or a SN impostor?
This analysis constrains the 2012b re-brightening of SN 2009ip to be a low-energy, asymmetric explosion in a complex medium. We interpret this 2012b episode to be the explosive ejection of the envelope of the massive progenitor star that later interacted with shells of material ejected during previous eruption episodes.
To unravel the nature of the physical mechanism behind the complex phenomenology of SN 2009ip, two key observational findings stand out: (i) its extreme similarity with SN 2010mc both in terms of time scales and energetics of the precursor bump and main explosion; (ii) the presence of a dominant time scale of ∼ 40 days, which regulates both the progenitor outburst history and major explosion (and which is also shared by completely independent events like SN 2010mc). While it is clear that the physical process leading to the 2012a eruption must have been fairly deepseated to unbind ∼ 0.1 M , it is unclear if any of the proposed mechanisms for substantial episodic mass loss (i.e. pulsational pair instability, super-Eddington fusion luminosity, Eddington-limit instabilities and shock heating of the stellar envelope) would be able to fully account for the observed properties, and in particular for the presence of dominant time scales. Indeed, the presence of dominant time scales might suggest a binary progenitor.
Finally, the extreme similarity to SN 2010mc allows us to conclude that the mechanism behind the precursor plus major outburst is not unique to SN 2009ip and is likely driven by few physical parameters. Moreover, this eruptive dynamic might represent an important evolutionary channel for massive stars. Future observations will reveal if SN 2009ip was able to survive.
We are grateful to E. Dwek, E. Quataert, H. Krimm, S. Barthelmy, I. Czekala, G. Chincarini for many interesting discussions and helpful suggestions. We would also like to thank the entire Swift team for their hard work and excellent support in scheduling the observations.
C.C.C. was supported at NRL by a Karles’ Fellowship and NASA DPR S-15633-Y. J.V. and T.S. are supported by the Hungarian OTKA Grant NN 107637. C.G. is supported by the NASA Postdoctoral Program (NPP). M.I. and C.C. were supported by the Creative Research Initiative program of the Korea Research Foundation (KRF) grant No. 2010-000712. R.C. acknowledges support from the National Science Foundation under grant AST-0807727 M. S. gratefully acknowledges generous support provided by the Danish Agency for Science and Technology and Innovation realized through a Sapere Aude Level 2 grant.

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The Dark Cosmology Centre is funded by the Danish National Science Foundation. The research of JCW, the Texas Supernova Group and EC is supported in part by NSF AST1109801 and by StScI grant HST-AR-12820. E.C. wishes to thank the University of Texas Graduate School for the William C. Powers fellowship given in support of his studies. The work of the Carnegie Supernova Project is supported by the National Science Foundation under grant AST1008343.
Based on observations made with ESO Telescopes at the La Silla Paranal Observatory under programme ID 090.D-0719. Observations reported here were obtained at the MMT Observatory, a joint facility of the Smithsonian Institution and the University of Arizona. This paper includes data gathered with the 6.5 meter Magellan Telescopes located at Las Campanas Observatory, Chile.
Observations were obtained with the JVLA operated by the National Radio Astronomy Observatory. The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.
Support for CARMA construction was derived from the Gordon and Betty Moore Foundation, the Kenneth T. and Eileen L. Norris Foundation, the James S. McDonnell Foundation, the Associates of the California Institute of Technology, the University of Chicago, the states of California, Illinois, and Maryland, and the National Science Founda-

tion. Ongoing CARMA development and operations are supported by the National Science Foundation under a cooperative agreement, and by the CARMA partner universities.
The Fermi LAT Collaboration acknowledges generous ongoing support from a number of agencies and institutes that have supported both the development and the operation of the LAT as well as scientific data analysis. These include the National Aeronautics and Space Administration and the Department of Energy in the United States, the Commissariat à l’Energie Atomique and the Centre National de la Recherche Scientifique / Institut National de Physique Nucléaire et de Physique des Particules in France, the Agenzia Spaziale Italiana and the Istituto Nazionale di Fisica Nucleare in Italy, the Ministry of Education, Culture, Sports, Science and Technology (MEXT), High Energy Accelerator Research Organization (KEK) and Japan Aerospace Exploration Agency (JAXA) in Japan, and the K. A. Wallenberg Foundation, the Swedish Research Council and the Swedish National Space Board in Sweden. Additional support for science analysis during the operations phase is gratefully acknowledged from the Istituto Nazionale di Astrofisica in Italy and the Centre National
d’Études Spatiales in France. This paper made use of the SUSPECT database
(http://www.nhn.ou.edu/ suspect/).

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APPENDIX
SHOCK BREAK-OUT IN A DENSE WIND SHELL: DIRECT CONSTRAINTS FROM OBSERVABLES
We provide here the set of equations we used to constrain the ejecta mass Mej, energy of the explosion E, the environment density ρw and mass-loss rate M˙ , starting from three observables: the radiated energy at break-out Erad, the bolometric light-curve rise-time trise and the radius at shock break-out Rbo. This work is based on Chevalier & Irwin (2011) . We consider here their model (a) (see their Fig. 1), where the break-out happens inside the dense wind shell of radius Rw so that Rbo < Rw. For SN 2009ip it turns out that Rbo Rw which leads to comparable estimates of the explosion and the environment parameters even if one were to use their model (b). The density in the wind shell is ρw = M˙ /(4πr2vw), or ρw = Dr−2 with D ≡ 5.0 × 1016(g/cm)D∗ and r is

38 Margutti et al.

in cgs units. We follow the parametrization by Chevalier & Irwin (2011) and solve their Eq. 1, Eq. 3 and Eq. 5 for the three observables. The system of equation is degenerate for Mej/E2. We obtain:

Mej = 10

Rbo 4.0 × 1014cm

−2

Erad 0.44 × 1050erg

−1

trise 6.6 days

2

E 1051erg

2
M

(A1)

D∗ =

trise 2 6.6 days

Rbo −3 4.0 × 1014cm

Erad 0.44 × 1050erg

M˙ = D∗

vw 1000 km s−1

M

yr−1

(A2) (A3)

k = 0.34

Rbo 4.0 × 1014cm

3

Erad 0.44 × 1050erg

−1

trise 6.6 days

−1cm2/g

(A4)

where vw is the wind velocity and k is the opacity. The mass of the wind shell enclosed within a radius R is:

Mw(r < R) = 6.3 × 1017

R cm

trise 6.6 days

2

Rbo 4.0 × 1014cm

−3

Erad 0.44 × 1050erg

g

(A5)

The luminosity from continued shock -wind interaction (Chevalier & Irwin 2011, their Eq. 9) after break-out can be written as:

L(t) = 7.1 × 1043

Erad 0.44 × 1050erg

trise −0.4

t

−0.6
erg

s−1

6.6 days

10 days

(A6)

The radius of the wind shell is:

Rw = 8.3 × 1013

trise 6.6 days

−0.8

Rbo 4.0 × 1014cm

tw cm days

(A7)

where tw is the time when the shock reaches the edge of the shell. In this scenario, the onset of the explosion is at:

t0 ≈ tpeak − 2trise

(A8)

Assuming hydrogen rich material with µp = 1, the column density from radius R to the observer can be written as (see Ofek

et al. 2013a):

N(R) ≈ 3.02 × 1025

trise 6.6 days

2

Rbo 4.0 × 1014cm

−3

Erad 0.44 × 1050erg

R 1015 cm

−1cm−2

(A9)

The bound-free optical depth at radius R as a function of the energy of the photons, E is τbf = Nσ(E) or

τbf(E) ≈ 3 × 103

trise 6.6 days

2

Rbo 4.0 × 1014cm

−3

Erad 0.44 × 1050erg

R −1 E −2.5 1015 cm keV

(A10)

adopting the approximated cross-section as in Ofek et al. (2013a). The free-free optical depth τff at radius R as a function of photon frequency ν is:

τff(ν) ≈ 2.6 × 108

Te 104 K

−1.35

ν GHz

−2.1

R 1015 cm

−3

trise 6.6 days

4

Rbo 4.0 × 1014cm

−6

Erad 0.44 × 1050erg

2

(A11)

where Te is the electron temperature.

TEMPORAL ANALYSIS
We identify the presence of a dominant time-scale of ∼ 40 days by applying the the Fast χ2 algorithm67 developed by Palmer (2009) to the R-band photometry obtained for SN 2009ip by Pastorello et al. (2012) in the period August 2009 - April 2012. This algorithm is suitable for irregularly sampled data with nonuniform errors and it is designed to identify the presence of periodic signals by minimizing the χ2 between the data and the model. We adopt a linear model for the trend. The best period was found to be T0 = 115 days. However, the reduction in the χ2 obtained for this period is comparable over a broad range of period candidates, thus proving that there it is no periodic or quasi–periodic signal. Assuming the Fourier decomposition with the fundamental period T0, among the first seven harmonics most (52%) of the temporal power is carried by the third harmonic (T0/3 = 38 days), followed by the fundamental (19%), the fifth harmonic (T0/5 = 23 days, 10%), the second and fourth harmonics (T0/2 = 58 and T0/4 = 29 days, 5% each). This is evidence for a dominant time-scale of about 40 days, with significant power distributed on time-scales between 30 and 50 days. Replicating the same analysis on time series obtained by randomizing the observed magnitudes has the effect of washing out the excess of power on those time scales. This demonstrates that the excess of power around 40 days is not due to the data sampling.
Additionally and independently, we calculated the Lomb–Scargle (LS; Lomb 1976; Scargle 1982) periodogram of the same time series. This is particularly suitable to unevenly sampled data sets and measures the power contributed by the different

67 http://public.lanl.gov/palmer/fastchi.html

SN2009ip

39

18 R pre-rebright

16 mean

1σ 14 90%

12

99% 99.9%

10

LS power

8

6

4

2

0

10

100

1000

10000

Period [d]

FIG. 32.— Lomb-Scargle periodogram of the R-band magnitude time series for the time interval August 2009 - April 2012 (i.e. before the 2012a and 2012b explosions). This plot shows significant power on time-scales around ∼ 40 days. We emphasize that this is not a claim for periodicity. Contours at various
significance levels are also shown.

TABLE 2 LOG OF OBSERVED Swift-UVOT SPECTRA

Date (UT)
2012-09-27 2012-09-28 2012-09-29 2012-09-30 2012-10-01 2012-10-02 2012-10-03 2012-10-04 2012-10-05 2012-10-06 2012-10-07 2012-10-08 2012-10-10 2012-10-12 2012-10-14 2012-10-16 2012-10-20 2012-10-22 2012-10-24 2012-10-26 2012-10-28 2012-11-02

Time range (UT)
09:03:41-14:37:27 13:41:01-16:53:23 10:27:52-12:24:46 05:48:22-17:19:31 01:13:23-19:11:08 14:08:21-22:05:35 02:49:01-15:52:51 09:10:01-14:19:17 09:12:55-15:33:06 04:26:57-11:09:55 07:41:54-08:05:27 07:44:55-12:57:21 03:01:55-12:56:01 06:20:54-11:31:19 04:50:54-06:54:14 06:32:56-11:35:19 08:30:16-13:41:40 09:59:14-13:30:15 13:19:14-17:00:26 10:10:14-13:53:54 08:41:21-09:00:03 17:09:16-19:03:52

Obs ID
31486008 31486011 31486012 32570001 32570002 31486020 32570003 32570004 32570005 32570006 32570007 32579001 32579002 32579003 32579004 32579005 32579007 32579008 32605001 32605002 32605003 32605004

Exposure (ks)
5.1 3.4 3.3 4.9 4.0 5.3 4.4 5.1 7.2 6.7 6.2 6.2 6.5 5.7 5.8 5.7 5.3 3.5 4.3 3.7 1.1 11.3

Roll Angle
212.0 215.0 215.0 216.0 216.0 229.0 220.0 220.0 220.0 220.0 220.0 243.0 243.0 243.0 243.0 243.0 243.0 243.0 249.0 249.0 249.0 249.0

frequencies to the total variance. This reduces to the Fourier power spectrum in the uniform sampling limit. As shown in Fig. 32, most power is concentrated between ∼ 30 and ∼50 days, with a peak around 40 days. To evaluate the impact of the aliasing and to quantify how significant the power is with respect to the white noise case (i.e. no preferred time scale), we carried out the following Monte Carlo simulation. We randomized the magnitudes among the same observation times and generated 104 synthetic profiles with the same variance as the real one. For each period we derived the corresponding confidence levels of 1σ, 90%, 99%, and 99.9%, as shown in Fig. 32. Comparing the real LS periodogram with the MC generated confidence levels, we conclude that power in excess of white noise around 40 days is significant at 99.9% confidence and we identify this time scale as the dominant one in the overall variance.
In addition, we applied both methods to the time profiles obtained for the V, H and R filters around the 2012 outburst. We applied both a third and a fourth degree polynomial de-trending. In all cases, we obtained similar results with dominant time scales in the range 30–50 days. This is a robust result which does not crucially depend on the polynomial degree used for de-trending the series.
OBSERVING LOGS
PHOTOMETRY TABLES

40 Margutti et al.

TABLE 3 LOG OF OBSERVED OPTICAL SPECTRA

Date (UT)
Aug 26.83 Sep 26.75 Sep 27.99
Sep 28.99 Sep 30.02 Oct 11.15 Oct 11.93 Oct 12.16 Oct 12.94 Oct 13.16 Oct 13.18 Oct 14.15 Oct 14.17 Oct 14.18 Oct 14.21 Oct 14.21 Oct 15.21 Oct 15.14 Oct 15.14 Oct 16.19 Oct 17.19 Oct 20.25 Oct 21.25 Oct 22.18 Oct 23.90 Oct 27.89 Oct 31.13 Nov 10.85 Nov 14.06 Nov 14.08 Nov 14.11 Nov 17.04 Nov 23.81 Dec 05.01 Dec 14.04 Dec 21.07 Jan 12.04 Apr 11.39

Telescope/Inst.
SALT/RSS SALT/RSS Magellan/MagEa Magellan/MagEb VLT/X-shooter KPNO/RCSpec SALT/RSS KPNO/RCSpec SALT/RSS KPNO/RCSpec FLWO/FAST KPNO/RCSpec MMT/Blue Channel MMT/Blue Channel MMT/Blue Channel FLWO/FAST FLWO/FAST KPNO/RCSpec MMT/Blue Channel FLWO/FAST FLWO/FAST FLWO/FAST FLWO/FAST FLWO/FAST SALT/RSS SALT/RSS VLT/X-shooter SALT/RSS MMT/Blue Channel MMT/Blue Channel MMT/Blue Channel Magellan/IMACS SALT/RSS Magellan/IMACS SOAR/GHTS MMT/Blue Channel Magellan/LDSS3 Magellan/IMACS

Range (Å)
3500−10000 3500−10000 3150.0−9400.0
3150−9400 3000−25000 3080−8760 3500−10000 3080−8760 3500−10000 5380−8290 3470−7414 5380−8290 3350−8570 5860−7160 3800−5120 3469−7413 3473−7417 5380−8290 3350−8570 3472−7416 3474−7418 3481−7422 3475−7419 3474−7418 3200-9000 3200-9000 3000−25000 3300−10500 3320−8540 5840−7140 3780−5110 4200 - 9300 3200 - 9000 3500 - 9400 3930 - 7985 3300 - 8500 3780−10500 4000−10200

R (λ/∆λ)
300 300 3400
3400 5100−8800
1200 300 1200 300 2500 2700 2500 1200 5000 3000 2700 2700 2500 1200 2700 2700 2700 2700 2700 1500 1500 5100−8800 300 1200 5000 3000 1300 1500 1300 300 300 650 1600

Grating (l/mm)
300 300 175
175 –
316 300 316 300 632 300 632 300 1200 1200 300 300 632 300 300 300 300 300 300 900 900
– 300 300 1200 1200 300 900 300 1390 740 VPH-all 300

Aperture (arcsec)
1.25 1.25 1.5
1.5 1.0/0.9/0.9c
1.5 1.25 1.5 1.25 1.5 3.0 1.5 1.0 1.0 1.0 3.0 3.0 1.5 1.0 3.0 3.0 3.0 3.0 3.0 1.25 1.25 1.0/0.9/0.9c 1.25 1.0 1.0 1.0 0.7 1.25 0.7 0.84 1.0 1.0 0.9

Airmass
2.84 1.16 1.25
1.23 1.10 2.15 1.24 2.09 1.29 2.07 2.04 2.10 2.03 2.03 2.12 2.07 2.08 2.11 2.03 2.05 2.04 2.20 2.80 2.04 1.20 1.21 1.23 1.25 2.11 2.06 2.04 1.02 1.25 1.25 1.44 2.46 2.75 1.87

Exp. Time (s)
900 1200 600
1200 744/556/800c
1200 300 1200 300 6000 1800 3600 300 600 900 1800 1800 4800 300 1500 1602 1500 1500 1500 600 600 744/556/800c 600 900 1000 1500 2700 600 27000 1800 2000 1200 1200

aObservations were obtained under poor 2 seeing conditions. b1.5 seeing. cRefers to UBV, VIS and NIR ranges, respectively

TABLE 4 LOG OF OBSERVED NIR SPECTRA

Date (UT)
Sep 27.21 Sep 29.19 Sep 30.20 Nov 05.11 Nov 19.09 Nov 25.11 Dec 03.04

Inst.
MDM MDM MDM FIRE FIRE FIRE FIRE

Range (Å)
9700 − 18000 9700 − 24300 9700 − 24300 8000 − 27400 8200 − 25000 8000 − 27200 8000 − 27200

R (λ/∆λ)
720 720 720 500 6000 500 500

Grating (l/mm)
··· ··· ··· ···
··· ···

Aperture (arcsec)
0.7 0.7 0.7 0.6 0.6 0.6 0.6

Airmass
2.00 2.00 2.00 1.24 1.28 1.67 1.31

Exp. Time (s)
1600 1600 1600 761 1200 1014 1014

SN2009ip

41

TABLE 5 Swift-UVOT PHOTOMETRY

Date (d)
84.68a 1174.88 1176.64 1183.44 1192.77 1196.24 1196.31 1196.37 1196.45 1198.48 1199.66 1200.99 1201.99 1202.26 1203.06 1204.05 1206.13 1208.67 1210.06 1212.20 1214.14 1216.21 1218.83 1220.55 1222.29 1224.49 1226.23 1228.71 1230.38 1232.86 1233.98 1234.98 1236.21 1238.09 1240.06 1242.29 1244.43 1246.91 1248.05 1250.05 1252.59 1254.09 1256.70 1258.53 1260.34 1262.21 1264.61 1266.15 1268.66 1270.35 1272.26 1274.56 1276.23 1278.40 1280.31 1282.15 1284.24 1286.08 1290.69 1294.47 1387.30

v (mag)
20.50(0.30) 17.10(0.21) 17.52(0.15) 17.74(0.15) 18.16(0.10) 14.91(0.05) 14.95(0.05) 14.86(0.05) 14.83(0.05) 14.38(0.04) 14.26(0.05) 14.14(0.05) 14.07(0.05) 14.11(0.05) 14.10(0.05) 13.91(0.05) 13.88(0.05) 13.92(0.05) 13.88(0.05) 13.95(0.05) 13.93(0.05) 14.13(0.05) 14.30(0.05) 14.30(0.05) 14.43(0.05) 14.46(0.05) 14.75(0.06) 15.01(0.06) 15.23(0.07) 15.09(0.06) 14.85(0.09) 14.86(0.09) 14.95(0.05) 15.16(0.05) 15.34(0.06) 15.68(0.06) 15.88(0.06) 16.00(0.06) 15.83(0.06) 16.04(0.07) 16.17(0.07) 16.33(0.06) 16.32(0.09) 16.42(0.07) 16.45(0.07) 16.41(0.06) 16.60(0.07) 16.71(0.07) 17.00(0.08) 17.28(0.10) 17.41(0.11) 17.94(0.14) 18.18(0.17) 18.29(0.27) 18.46(0.17) 18.51(0.20) 18.70(0.20) 18.67(0.25) 18.72(0.18) 18.84(0.22) 20.18(0.42)

Date (d)
84.68 1174.90 1176.64 1183.44 1190.75 1192.76 1196.23 1196.30 1196.36 1196.45 1197.70 1198.48 1199.65 1200.98 1201.98 1202.25 1203.06 1203.32 1204.05 1204.52 1205.52 1206.12 1208.67 1210.06 1212.20 1214.14 1216.27 1218.83 1220.55 1222.29 1224.49 1226.23 1228.71 1230.37 1232.85 1234.48 1236.18 1238.08 1240.05 1242.29 1244.46 1246.91 1248.04 1250.01 1252.59 1254.09 1256.66 1258.53 1260.30 1262.20 1264.61 1266.14 1268.66 1270.35 1272.25 1274.55 1276.23 1279.40 1283.17 1287.48 1291.36 1294.46 1387.30

b (mag)
20.76(0.19) 17.29(0.05) 17.24(0.07) 17.62(0.08) 18.39(0.15) 18.13(0.06) 14.81(0.05) 14.75(0.05) 14.77(0.05) 14.74(0.05) 14.39(0.05) 14.28(0.04) 14.11(0.05) 13.94(0.05) 13.98(0.05) 13.96(0.05) 13.96(0.05) 13.96(0.05) 13.84(0.05) 13.76(0.13) 13.73(0.11) 13.75(0.05) 13.86(0.05) 13.91(0.05) 13.91(0.05) 13.97(0.05) 14.18(0.07) 14.39(0.05) 14.44(0.05) 14.57(0.05) 14.68(0.05) 14.94(0.05) 15.25(0.06) 15.50(0.06) 15.38(0.06) 15.16(0.06) 15.25(0.06) 15.40(0.05) 15.73(0.05) 16.07(0.05) 16.23(0.08) 16.42(0.06) 16.33(0.06) 16.47(0.07) 16.80(0.06) 16.85(0.06) 17.00(0.06) 17.16(0.06) 17.20(0.08) 17.11(0.06) 17.34(0.07) 17.55(0.07) 18.04(0.09) 18.00(0.09) 18.35(0.10) 18.95(0.16) 19.05(0.17) 19.35(0.15) 19.15(0.13) 19.48(0.18) 19.44(0.20) 19.61(0.20) 20.32(0.20)

Date (d)
84.68 1174.90 1176.64 1183.44 1190.75 1192.76 1196.23 1196.30 1196.36 1196.45 1197.70 1198.47 1199.65 1200.98 1201.98 1200.47 1202.25 1203.06 1203.12 1203.32 1203.39 1204.05 1204.38 1205.38 1205.52 1205.65 1206.12 1206.25 1208.39 1208.52 1208.52 1208.67 1210.06 1210.19 1210.46 1212.20 1212.33 1212.33 1212.46 1214.14 1216.20 1216.27 1216.27 1216.40 1216.41 1218.82 1220.55 1220.45 1222.29 1224.62 1224.49 1226.23 1228.71 1228.36 1230.37 1232.85 1233.98 1234.98 1236.18 1238.08 1240.05 1242.29 1244.43 1246.91 1248.04 1250.04 1252.59 1254.09 1256.66 1258.53 1260.33 1262.20 1264.61 1266.14 1268.66 1387.30

u (mag)
20.08(0.17) 16.48(0.06) 16.46(0.07) 17.23(0.09) 17.77(0.12) 17.88(0.07) 13.46(0.05) 13.44(0.05) 13.39(0.05) 13.36(0.05) 12.99(0.05) 12.90(0.04) 12.78(0.04) 12.65(0.04) 12.61(0.04) 12.72(0.04) 12.64(0.04) 12.63(0.04) 12.67(0.04) 12.63(0.06) 12.67(0.06) 12.60(0.04) 12.56(0.04) 12.57(0.04) 12.54(0.04) 12.56(0.04) 12.50(0.04) 12.51(0.04) 12.63(0.04) 12.60(0.06) 12.61(0.04) 12.63(0.04) 12.80(0.04) 12.75(0.04) 12.75(0.04) 12.76(0.04) 12.78(0.06) 12.80(0.04) 12.81(0.04) 12.84(0.04) 13.09(0.04) 13.05(0.06) 13.07(0.04) 13.06(0.06) 13.09(0.04) 13.48(0.04) 13.66(0.04) 13.69(0.04) 13.75(0.04) 14.02(0.04) 14.02(0.04) 14.29(0.04) 14.85(0.05) 14.81(0.04) 15.04(0.05) 14.99(0.05) 14.90(0.08) 14.77(0.07) 14.87(0.05) 15.06(0.05) 15.43(0.05) 15.82(0.05) 16.13(0.06) 16.23(0.06) 16.18(0.06) 16.44(0.07) 16.75(0.07) 17.07(0.07) 17.30(0.09) 17.24(0.09) 17.30(0.08) 17.26(0.08) 17.41(0.09) 17.67(0.10) 18.11(0.12) 20.88(0.38)

Date (d)
84.67 1174.89 1176.63 1183.43 1190.75 1192.76 1196.34 1197.47 1197.70 1198.47 1198.61 1199.47 1199.65 1200.98 1201.98 1202.25 1203.06 1204.05 1206.12 1208.67 1210.06 1212.20 1214.13 1218.82 1220.55 1222.28 1224.49 1226.23 1228.70 1230.37 1232.85 1234.48 1236.21 1238.08 1240.09 1242.29 1244.46 1246.93 1248.04 1250.04 1252.59 1254.09 1256.66 1258.52 1260.33 1262.20 1264.60 1266.14 1268.66 1270.34 1272.25 1274.55 1276.22 1278.39 1280.30 1282.14 1285.14 1290.68 1294.46 1387.30

w1 (mag)
20.48(0.19) 17.09(0.07) 16.96(0.08) 18.08(0.11) 19.08(0.21) 19.23(0.16) 12.98(0.04) 12.75(0.04) 12.62(0.04) 12.51(0.04) 12.59(0.04) 12.53(0.04) 12.41(0.04) 12.30(0.04) 12.31(0.04) 12.31(0.04) 12.30(0.04) 12.22(0.04) 12.17(0.04) 12.44(0.04) 12.55(0.04) 12.69(0.04) 12.77(0.04) 13.71(0.04) 13.95(0.04) 14.18(0.04) 14.40(0.04) 14.75(0.05) 15.30(0.05) 15.48(0.06) 15.64(0.06) 15.52(0.06) 15.53(0.05) 15.72(0.05) 16.01(0.05) 16.27(0.05) 16.63(0.09) 16.94(0.08) 16.88(0.07) 17.05(0.08) 17.34(0.08) 17.50(0.07) 17.83(0.09) 17.93(0.09) 18.04(0.09) 18.42(0.11) 18.32(0.11) 18.41(0.11) 18.73(0.13) 18.70(0.14) 18.85(0.14) 19.05(0.16) 19.16(0.17) 19.42(0.30) 19.64(0.21) 19.78(0.27) 19.78(0.19) 20.10(0.24) 19.66(0.18) 21.12(0.33)

Date (d)
84.68 1174.90 1176.64 1183.44 1192.77 1196.34 1197.46 1197.71 1198.48 1198.59 1199.65 1200.99 1201.99 1201.38 1202.26 1203.06 1200.49 1201.43 1203.27 1204.05 1204.50 1205.53 1206.12 1206.33 1208.67 1207.44 1208.44 1210.34 1210.06 1212.20 1214.14 1216.21 1216.47 1218.83 1212.38 1214.25 1216.39 1220.47 1222.50 1220.55 1222.29 1224.49 1226.23 1224.64 1226.44 1228.71 1230.37 1232.86 1233.79 1228.38 1234.48 1234.26 1236.21 1238.08 1240.05 1242.29 1244.43 1246.94 1248.05 1250.04 1252.59 1254.09 1256.66 1258.53 1260.34 1262.20 1264.61 1266.14 1268.66 1271.32 1275.19 1280.34 1285.14 1290.69 1294.46 1387.30

w2 (mag)
20.77(0.17) 17.85(0.09) 17.86(0.10) 19.09(0.14) 20.31(0.18) 12.83(0.05) 12.66(0.05) 12.52(0.05) 12.42(0.04) 12.43(0.04) 12.32(0.04) 12.20(0.04) 12.19(0.04) 12.23(0.04) 12.21(0.04) 12.23(0.04) 12.26(0.04) 12.23(0.04) 12.27(0.04) 12.16(0.04) 12.13(0.04) 12.13(0.04) 12.14(0.04) 12.18(0.04) 12.47(0.04) 12.30(0.04) 12.45(0.04) 12.70(0.04) 12.65(0.04) 12.80(0.04) 12.94(0.04) 13.44(0.04) 13.50(0.04) 14.22(0.04) 12.82(0.04) 12.97(0.04) 13.47(0.04) 14.59(0.04) 14.84(0.04) 14.58(0.04) 14.86(0.04) 15.19(0.05) 15.54(0.06) 15.21(0.04) 15.56(0.06) 16.10(0.06) 16.32(0.06) 16.55(0.06) 16.48(0.05) 15.94(0.07) 16.52(0.07) 16.50(0.04) 16.51(0.05) 16.70(0.05) 17.00(0.06) 17.36(0.06) 17.71(0.07) 17.82(0.10) 17.87(0.08) 18.18(0.09) 18.49(0.10) 18.72(0.10) 18.86(0.14) 19.01(0.12) 19.27(0.13) 19.17(0.12) 19.27(0.14) 19.46(0.15) 19.45(0.15) 19.82(0.17) 20.11(0.18) 20.33(0.18) 20.40(0.20) 20.55(0.21) 20.54(0.22) 21.56(0.30)

Date (d)
84.69 1176.64 1183.44 1192.74 1196.35 1198.48 1200.99 1201.99 1202.75 1202.26 1203.07 1204.06 1206.13 1208.68 1210.07 1212.20 1214.14 1216.21 1216.41 1218.83 1220.55 1222.29 1222.41 1224.49 1226.49 1226.23 1228.71 1230.38 1232.86 1234.41 1234.98 1236.18 1238.09 1240.06 1242.30 1242.30 1244.43 1246.94 1248.05 1250.05 1252.59 1254.09 1256.70 1258.54 1260.34 1262.21 1264.61 1266.15 1268.66 1271.32 1275.20 1280.35 1286.55 1292.40 1387.30

m2 (mag)
20.87(0.23) 17.73(0.09) 18.92(0.12) 20.25(0.22) 12.78(0.04) 12.30(0.04) 12.10(0.04) 12.12(0.04) 12.17(0.04) 12.14(0.04) 12.13(0.04) 12.06(0.04) 12.04(0.04) 12.30(0.04) 12.47(0.04) 12.58(0.04) 12.71(0.04) 13.16(0.04) 13.20(0.06) 13.88(0.04) 14.22(0.04) 14.49(0.04) 14.49(0.05) 14.82(0.04) 15.20(0.05) 15.16(0.05) 15.64(0.05) 15.92(0.06) 16.21(0.06) 16.21(0.05) 16.18(0.08) 16.13(0.06) 16.31(0.05) 16.58(0.06) 16.93(0.06) 16.93(0.06) 17.28(0.07) 17.39(0.08) 17.57(0.07) 17.85(0.08) 18.32(0.10) 18.32(0.08) 18.57(0.13) 18.78(0.10) 18.65(0.09) 18.95(0.11) 18.99(0.12) 19.28(0.13) 19.43(0.14) 19.77(0.15) 20.00(0.16) 19.84(0.18) 20.00(0.15) 20.31(0.16)
>21.90

aDates are in MJD-55000 (days).

42 Margutti et al.

TABLE 6 R AND I BAND PHOTOMETRY

Date (MJD)
56181.25 56194.25 56199.50 56201.50 56204.50 56205.50 56206.50 56209.00 56209.50 56210.50 56215.50 56218.50 56219.50 56220.50 56221.50 56223.50 56224.50 56226.50 56226.50 56227.50 56228.50 56234.50 56238.50 56238.50 56239.50 56240.50 56240.50 56242.50 56245.50 56246.50 56247.50 56247.50 56248.00 56248.50 56250.50

R (mag)
16.54(0.04) 17.08(0.06) 14.01(0.01) 13.86(0.01) 13.70(0.01) 13.66(0.01) 13.65(0.01) 13.70(0.04) 13.70(0.01) 13.71(0.01) 13.75(0.01) 13.95(0.01) 14.01(0.01) 14.04(0.01) 14.04(0.03) 14.17(0.02) 14.22(0.01) 14.43(0.01) 14.40(0.03) 14.54(0.02) 14.65(0.02) 14.59(0.04) 14.78(0.05) 14.78(0.01) 14.87(0.02) 14.98(0.04) 14.98(0.02) 15.15(0.02) 15.34(0.02) 15.36(0.03) 15.36(0.08) 15.37(0.02) 15.37(0.04) 15.38(0.03) 15.68(0.06)

Date (MJD)
56181.25 56194.25 56197.00 56198.00 56199.25 56200.00 56201.00 56202.00 56203.00 56204.00 56205.00 56206.00 56208.00 56209.00 56209.25 56210.00 56211.00 56212.00 56213.00 56214.00 56215.25 56216.00 56217.00 56218.00 56219.00 56220.00 56221.00 56222.00 56223.00 56224.00 56225.00 56226.00 56229.00 56230.00 56231.00 56233.00 56234.00 56235.00 56236.00 56237.00 56238.00 56239.00 56240.00 56241.00 56242.00 56243.00 56244.00 56245.00 56246.00 56247.00 56248.00 56248.00 56249.00 56250.00 56251.00

I (mag)
16.55(0.07) 17.82(0.09) 14.43(0.13) 14.17(0.07) 13.94(0.06) 13.97(0.06) 13.89(0.04) 13.79(0.08) 13.86(0.08) 13.77(0.07) 13.69(0.07) 13.65(0.05) 13.67(0.07) 13.66(0.07) 13.60(0.09) 13.69(0.07) 13.59(0.07) 13.65(0.06) 13.56(0.10) 13.56(0.07) 13.69(0.09) 13.63(0.08) 13.74(0.09) 13.89(0.08) 13.83(0.06) 13.83(0.08) 13.90(0.10) 14.00(0.07) 14.04(0.13) 14.16(0.10) 14.14(0.06) 14.19(0.07) 14.44(0.12) 14.77(0.09) 14.55(0.09) 14.52(0.10) 14.44(0.08) 14.53(0.08) 14.36(0.06) 14.43(0.07) 14.60(0.12) 14.74(0.08) 14.70(0.08) 14.90(0.07) 14.98(0.14) 14.97(0.19) 14.99(0.08) 15.14(0.14) 15.26(0.08) 15.19(0.08) 15.20(0.06) 15.10(0.10) 15.00(0.09) 15.26(0.09) 15.19(0.09)

SN2009ip

TABLE 7 NIR PHOTOMETRY FROM PAIRITEL

Date (MJD)
56213.19 56214.19 56215.20 56216.21 56219.18 56223.21 56225.14 56226.14 56227.15 56228.15 56230.17 56231.17 56232.17 56236.11 56237.11 56238.12 56243.10 56245.13 56246.14 56248.09 56253.11 56255.10 56256.07 56257.08 56266.07 56267.08

J (mag)
13.53(0.01) 13.60(0.01) 13.56(0.01) 13.61(0.02) 13.75(0.03) 13.88(0.01) 13.99(0.02) 14.08(0.02) 14.18(0.02) 14.33(0.02) 14.43(0.03) 14.41(0.02) 14.43(0.03) 14.26(0.03) 14.35(0.02) 14.34(0.02) 14.74(0.04) 14.80(0.03) 14.88(0.10) 14.88(0.04) 15.07(0.05) 15.03(0.05) 15.15(0.09) 15.15(0.05) 15.21(0.08) 15.22(0.06)

Date (MJD)
56213.19 56214.19 56215.20 56216.21 56219.18 56223.21 56225.14 56226.14 56227.15 56228.15 56230.17 56231.17 56232.17 56236.11 56237.11 56238.12 56245.13 56248.09 56253.11 56255.10 56256.07 56257.08 56266.07 56267.08

H (mag)
13.37(0.02) 13.45(0.03) 13.41(0.02) 13.51(0.04) 13.60(0.06) 13.85(0.03) 13.90(0.04) 13.90(0.03) 14.08(0.04) 14.22(0.05) 14.38(0.05) 14.32(0.05) 14.26(0.07) 14.11(0.05) 14.12(0.04) 14.22(0.05) 14.65(0.06) 14.66(0.11) 14.88(0.10) 14.82(0.14) 15.09(0.19) 14.96(0.12) 14.94(0.13) 15.27(0.15)

Date (MJD)
56213.19 56214.19 56215.20 56216.21 56219.18 56226.14 56227.15 56228.15 56230.17 56231.17 56232.17 56236.11 56237.11 56238.12 56243.10 56255.10 56266.07 56267.08

K (mag)
13.25(0.05) 13.20(0.06) 13.24(0.05) 13.23(0.07) 13.44(0.12) 13.76(0.10) 13.66(0.08) 14.02(0.12) 14.11(0.12) 14.19(0.17) 14.10(0.15) 14.01(0.15) 13.86(0.12) 13.71(0.16) 14.35(0.14) 14.93(0.29) 14.82(0.25) 14.53(0.36)

43