A Study of the Type Ia/IIn Supernova 2005gj from X-ray to the Infrared: Paper I
J.L. Prieto, P.M. Garnavich, M.M. Phillips, D.L. DePoy, J.Parrent, D. Pooley, V.V. Dwarkadas, E. Baron, B. Bassett, A. Becker, D. Cinabro, F. DeJongh, B. Dilday, M. Doi, J.A. Frieman, C.J. Hogan, J. Holtzman, S. Jha, R. Kessler, K. Konishi, H. Lampeitl, J. Marriner, J. L. Marshall, G. Miknaitis, R.C. Nichol, A.G. Riess, M.W. Richmond, R. Romani, M. Sako, D.P. Schneider, M. Smith, N. Takanashi, K. Tokita, K. van der Heyden, N. Yasuda, C. Zheng, J.C. Wheeler, J. Barentine, J. Dembicky, J. Eastman, S. Frank, W. Ketzeback, R.J. McMillan, N. Morrell, G. Folatelli, C. Contreras, C.R. Burns, W. L. Freedman, S. Gonzalez, M. Hamuy, W. Krzeminski, B.F. Madore, D. Murphy, S.E. Persson, M. Roth, N.B. Suntzeff
aa r X i v : . [ a s t r o - ph ] J un A Study of the Type Ia/IIn Supernova 2005gj from X-ray to theInfrared: Paper I , , , J. L. Prieto , P. M. Garnavich , M. M. Phillips , D. L. DePoy , J. Parrent , D. Pooley , ,V. V. Dwarkadas , E. Baron , , B. Bassett , , A. Becker , D. Cinabro , F. DeJongh ,B. Dilday , , M. Doi , J. A. Frieman , , , C. J. Hogan , J. Holtzman , S. Jha ,R. Kessler , , K. Konishi , H. Lampeitl , J. Marriner , J. L. Marshall , G. Miknaitis ,R. C. Nichol , A. G. Riess , , M. W. Richmond , R. Romani , M. Sako ,D. P. Schneider , M. Smith , N. Takanashi , K. Tokita , K. van der Heyden ,N. Yasuda , C. Zheng , J. C. Wheeler, , J. Barentine , , J. Dembicky , J. Eastman ,S. Frank , W. Ketzeback , R. J. McMillan , N. Morrell , G. Folatelli , C. Contreras ,C. R. Burns , W. L. Freedman , S. Gonz´alez , M. Hamuy , W. Krzeminski ,B. F. Madore , D. Murphy , S. E. Persson , M. Roth , N. B. Suntzeff Based in part on observations obtained with the Apache Point Observatory 3.5-meter telescope, whichis owned and operated by the Astrophysical Research Consortium. Based in part on observations taken at the Cerro Tololo Inter-American Observatory, National OpticalAstronomy Observatory, which is operated by the Association of Universities for Research in Astronomy,Inc. (AURA) under cooperative agreement with the National Science Foundation. This paper includes data gathered with the 6.5 meter Magellan Telescopes located at Las CampanasObservatory, Chile. Partly based on observations collected at the European Southern Observatory, Chile, in the course ofprogramme 076.A-0156. Department of Astronomy, Ohio State University, 140 West 18th Avenue, Columbus, OH 43210-1173. University of Notre Dame, 225 Nieuwland Science, Notre Dame, IN 46556-5670. Las Campanas Observatory, Carnegie Observatories, La Serena, Chile Holmer L. Dodge Department of Physics and Astronomy, University of Oklahoma, 440 West Brooks,Room 100, Norman, OK 73019-2061 Astronomy Department, University of California at Berkeley, 601 Campbell Hall, Berkeley, CA 94720 Chandra Fellow Department of Astronomy and Astrophysics, The University of Chicago, 5640 South Ellis Avenue,Chicago, IL 60637. Computational Research Division, Lawrence Berkeley National Laboratory, MS 50F-1650, 1 CyclotronRd, Berkeley, CA 94720 USA Department of Mathematics and Applied Mathematics, University of Cape Town, Rondebosch 7701,South Africa. South African Astronomical Observatory, P.O. Box 9, Observatory 7935, South Africa. Department of Astronomy, University of Washington, Box 351580, Seattle, WA 98195. Department of Physics, Wayne State University, Detroit, MI 48202. Center for Particle Astrophysics, Fermi National Accelerator Laboratory, P.O. Box 500, Batavia, IL60510. Kavli Institute for Cosmological Physics, The University of Chicago, 5640 South Ellis Avenue Chicago,IL 60637. Department of Physics, University of Chicago, Chicago, IL 60637. Institute of Astronomy, Graduate School of Science, University of Tokyo 2-21-1, Osawa, Mitaka, Tokyo181-0015, Japan. Department of Astronomy, MSC 4500, New Mexico State University, P.O. Box 30001, Las Cruces, NM88003. [email protected]
ABSTRACT
We present extensive ugriz Y HJ K s photometry and optical spectroscopy ofSN 2005gj obtained by the SDSS-II and CSP Supernova Projects, which giveexcellent coverage during the first 150 days after the time of explosion. Thesedata show that SN 2005gj is the second clear case, after SN 2002ic, of a ther-monuclear explosion in a dense circumstellar environment. Both the presence ofsingly and doubly ionized iron-peak elements (Fe III and weak S II, Si II) nearmaximum light as well as the spectral evolution show that SN 2002ic-like eventsare Type Ia explosions. Independent evidence comes from the exponential decayin luminosity of SN 2005gj, pointing to an exponential density distribution ofthe ejecta. The interaction of the supernova ejecta with the dense circumstellar Kavli Institute for Particle Astrophysics & Cosmology, Stanford University, Stanford, CA 94305-4060. Enrico Fermi Institute, University of Chicago, 5640 South Ellis Avenue, Chicago, IL 60637. Institute for Cosmic Ray Research, University of Tokyo, 5-1-5, Kashiwanoha, Kashiwa, Chiba, 277-8582,Japan. Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218. Institute of Cosmology and Gravitation, Mercantile House, Hampshire Terrace, University ofPortsmouth, Portsmouth PO1 2EG, UK. Department of Physics and Astronomy, Johns Hopkins University, 3400 North Charles Street, Baltimore,MD 21218. Physics Department, Rochester Institute of Technology, 85 Lomb Memorial Drive, Rochester, NY14623-5603. Department of Physics and Astronomy, University of Pennsylvania, 203 South 33rd Street, Philadelphia,PA 19104. Department of Astronomy and Astrophysics, The Pennsylvania State University, 525 Davey Laboratory,University Park, PA 16802. Department of Astronomy, McDonald Observatory, University of Texas, Austin, TX 78712 Apache Point Observatory, P.O. Box 59, Sunspot, NM 88349. Carnegie Institution of Washington, 813 Santa Barbara St., Pasadena, CA 91101 Universidad de Chile, Departamento de Astronom´ıa, Santiago, Chile Texas A&M University Physics Department, College Station, TX α componentssuggest that the CSM around SN 2005gj is clumpy and it has a flatter density dis-tribution compared with the steady wind solution, in agreement with SN 2002ic.An early X-ray observation with Chandra gives an upper-limit on the mass lossrate from the companion of ˙ M . × − M ⊙ yr − . Subject headings: supernovae: general — supernovae: individual (SN 2005gj)
1. Introduction
Thermonuclear supernova explosions (Type Ia supernovae, SN Ia hereafter) are believedto be the detonation or deflagration of a white dwarf accreting matter from a companionstar (Arnett 1982). The mass of the white dwarf slowly increases until it approaches theChandrasekhar limit where the star becomes thermally unstable. At this point fusion ofCarbon and Oxygen begins near the center and quickly moves through most of the star beforedegeneracy is lifted. The result is a spectacular and powerful explosion that is visible acrossmuch of the Universe. Since SN Ia arise from a narrow range of white dwarf masses, their peakluminosities are very consistent and they make excellent distance indicators (e.g., Phillips1993). SN Ia are powerful probes of cosmology and have been instrumental in narrowing theuncertainty in the Hubble parameter, discovery of the accelerating universe, and constrainingdark energy models (Hamuy et al. 1995, 1996b; Riess et al. 2005, 1998; Perlmutter et al.1999; Riess et al. 2004; Astier et al. 2006; Riess et al. 2006; Wood-Vasey et al. 2007).But the use of SN Ia as reliable distance indicators will always be questioned until theprogenitor and explosion physics are well-understood. What types of binaries create SN Ia?How is matter transferred to the white dwarf without causing thermonuclear runaways onthe surface? Are there several types of progenitors?. These big questions remain to beanswered and detailed observations of hundreds of events have yielded few clues.In 2002, Hamuy et al. (2003) identified a new kind of supernova. The early spectrumof SN 2002ic was a cross between a Type Ia event and a Type IIn (Deng et al. 2004, havecalled this type of supernova a “IIa”), showing P-Cygni features similar to SN Ia and resolvedBalmer lines in emission. Type IIn supernovae are core-collapse explosions going off in densecircumstellar environments (Schlegel 1990; Chevalier & Fransson 1994). They are relativelycommon since the massive stars that create core-collapse supernovae often have thick winds.If the interpretations of the pre-explosion observations of SN 2005gl are correct, SN IIn could 5 –be associated in some cases with luminous blue variables (Gal-Yam et al. 2006).In the case of SN 2002ic, the presence of Balmer lines with profiles characteristic ofSN IIn, and the high luminosities and slow decline after maximum lead to the conclusionthat most of the energy came from the interaction of the ejecta with a dense circumstel-lar medium (CSM). Other Type IIn events (SN 1997cy, Germany et al. 2000; SN 1999E,Rigon et al. 2003) have been re-classified as SN 2002ic-like that were caught late in theirevolution (Hamuy et al. 2003; Wood-Vasey et al. 2004).SN 2002ic provided the first direct evidence that thermonuclear explosions can also oc-cur in a dense medium, but in this case the circumstellar medium is probably generated by anAsymptotic Giant companion (Hamuy et al. 2003; Wang et al. 2004; Han & Podsiadlowski2006). However, there is still debate in the literature about the origin of SN 2002ic.Livio & Riess (2003) proposed the merger of two white dwarfs as a possible progenitor,with the explosion occurring in the common envelope phase. Chugai et al. (2004) concludedthat the properties of the circumstellar interaction in 2002ic-like events can be broadly ex-plained by the SN 1.5 scenario (Iben & Renzini 1983): the thermonuclear explosion of thedegenerate core of a massive AGB star. Recently, Benetti et al. (2006) questioned the earlierinterpretation of the observations and proposed that SN 2002ic can be equally well explainedby the core collapse of a stripped-envelope massive star in a dense medium.SN 2005gj was discovered on 2005 September 28.6 (UT) by the SDSS-II Supernova Col-laboration (Frieman et al. 2007) in gri images obtained with the SDSS-2.5m telescope atApache Point Observatory (APO). The new supernova (Barentine et al. 2005) was ∼ ′′ fromthe center of its host galaxy at the position α = 03 h m . s δ = +00 ◦ ′ . ′′ g, r, i ) = (18 . , . , .
7) mag, obtained from PSF photometry afterkernel-matching and subtraction of a template image in each band. The SN was indepen-dently discovered by the Nearby Supernova Factory on 2005 September 29 (Aldering et al.2006).SN 2005gj was classified as a Type Ia candidate from the first three epochs of the gri light curve using a light curve fitting program, and was sent to the queue of the MDM-2.4mtelescope for spectroscopic confirmation. The optical spectrum obtained on 2005 October1 (UT) showed a blue continuum with resolved Hydrogen Balmer lines in emission, verysimilar to the spectrum of a young Type IIn supernova, but also with an unusual continuumshowing broad and weak absorption features. Further spectroscopic follow-up showed adramatic evolution. The continuum became substantially redder and developed broad, P-Cygni features probably associated with blended lines of Fe-peak mass elements, similar to aType Ia SN a few weeks after maximum. The spectrum obtained on 2005 Nov. 12 (UT) wasremarkably similar to that of the unusual Type Ia supernova SN 2002ic obtained on 2002 6 –Dec. 27 (UT) (Prieto et al. 2005).Aldering et al. (2006) presented optical photometry and spectroscopy of this SN. Throughdetailed analysis they confirmed its photometric and spectroscopic resemblance to SN 2002ic,confirming it was a new case of a Type Ia explosion interacting with a dense circumstellarenvironment. From a spectrum obtained with the slit oriented to overlap with its host galaxy,they calculated a redshift for the host of z = 0 . ± . Swift (Immler et al. 2005).Here we present extensive follow-up photometry and spectroscopy of the Type Ia/Type IInSN 2005gj during the first ∼
150 days after discovery. These are the most detailed observa-tions ever obtained of a SN 2002ic-like event, and provide insight into the early evolution,progenitor, and variety of these events. We also present a sensitive, early X-ray observationwith
Chandra that gives an upper limit on the X-ray luminosity of this peculiar object. Wedescribe the optical and NIR photometry of SN 2005gj in § §
3. We describe the X-ray observation with
Chandra in §
4. An analysis of the photo-metric and spectroscopic data are presented in §
5. Finally, we discuss the significance ofthe results in §
6. We will present later photometric follow-up and Spitzer/IRAC observa-tions in Paper II. We adopt a cosmology with H = 72 ± − Mpc − , Ω M = 0 .
3, andΩ Λ = 0 . µ = 37 .
15 mag to the host of SN 2005gj.
2. Photometry2.1. SDSS and MDM
The Sloan Digital Sky Survey (SDSS; York et al. 2000) uses a wide-field, 2.5-metertelescope (Gunn et al. 2006) and mosaic CCD camera (Gunn et al. 1998) at APO to sur-vey the sky. The SDSS-II Supernova Survey, part of a 3-year extension of the originalSDSS, uses the APO-2.5m telescope to detect and measure light-curves for a large num-ber of supernovae through repeat scans of the SDSS Southern equatorial stripe (about2.5 deg wide by ∼
120 deg long) over the course of three 3-month campaigns (Sept-Nov.2005-2007). SN 2005gj was discovered in the second month of the first campaign (Octo-ber 2005). Twenty epochs of ugriz photometry were obtained between 2005 Sep 26-Nov 30(U.T.). Details of the photometric system, magnitude system, astrometry and calibration aregiven in Fukugita et al. (1996), Lupton et al. (1999), Hogg et al. (2001), Smith et al. (2002),Pier et al. (2003), Ivezi´c et al. (2004), and Tucker et al. (2006). Additional griz imaging of 7 –SN 2005gj was obtained with the MDM Observatory 2.4m telescope using a facility CCDimager (RETROCAM; see Morgan et al. 2005, for a complete description of the imager).Photometry of SN 2005gj on the SDSS images was carried out using the scene modelingcode developed for SDSS-II as described in Holtzman et al. (2007). A sequence of starsaround the supernova were taken from the list of Ivezi´c et al. (2007), who derived standardSDSS magnitudes from multiple observations taken during the main SDSS survey underphotometric conditions. Using these stars, frame scalings and astrometric solutions werederived for each of the supernova frames, as well as for the twenty five pre-supernova framestaken as part of either the main SDSS survey or the SN survey. Finally, the entire stack offrames was simultaneously fit for a single supernova position, a fixed galaxy background ineach filter (characterized by a grid of galaxy intensities), and the supernova brightness ineach frame.The supernova photometry in the MDM frames was also determined using the scenemodeling code. Since the MDM observations had different response functions from thestandard SDSS bandpasses, the photometric frame solutions included color terms from theSDSS standard magnitudes. To prevent uncertainties in the frame parameters and colorterms from possibly corrupting the galaxy model (here affecting the SDSS photometry),the MDM data were not included in the galaxy determination, but the galaxy model asdetermined from the SDSS was used (with color terms) to subtract the galaxy in the MDMframes. The resulting SN photometry from the MDM frames is reported on the native MDMsystem, since the color terms derived from stars are likely not to apply to the spectrum ofthe supernova.Figure 1 shows a 3 . ′ × . ′ field around SN 2005gj and 16 comparison stars used forcalibration of the SN magnitudes by SDSS and the Carnegie Supernova Group (CSP; see § ugriz and CSP u ′ g ′ r ′ i ′ photometryof the comparison stars in common. The final SDSS and MDM griz photometry is given inTable 2. Optical photometry from the CSP was obtained with the Swope-1m telescope at LCO,using a SITe CCD and a set of u ′ g ′ r ′ i ′ filters. A subraster of 1200 × ′′ pixel − , yielded a field of viewof 8 . ′ × . ′ . Typical image quality ranged between 1 ′′ and 2 ′′ (FWHM). A photometricsequence of comparison stars in the SN field was calibrated with the Swope telescope from 8 –observations of SDSS standard stars from Smith et al. (2002) during four photometric nights.SN magnitudes in the u ′ g ′ r ′ i ′ system were obtained differentially relative to the comparisonstars using PSF photometry. In order to minimize the contamination from the host galaxylight in the SN magnitudes, PSF-matched ugri template images from SDSS were subtractedfrom the SN images. On every galaxy-subtracted image, a PSF was fitted to the SN andcomparison stars within a radius of 3 ′′ . See Hamuy et al. (2006) for further details aboutthe measurement procedures.The NIR photometry of SN 2005gj was obtained by the CSP using three different instru-ments/telescopes. A total of 15 epochs in Y , J and H filters were obtained using RetroCam,mounted on the Swope-1m telescope at LCO. A few additional epochs in Y J HK s were ob-tained with the Wide Infrared Camera (WIRC; Persson et al. 2002) mounted on the duPont-2.5m telescope, and the PANIC camera (Martini et al. 2004) mounted on the Magellan-6.5mBaade telescope, both at LCO. We refer to Hamuy et al. (2006) and Phillips et al. (2007)for details of the imagers and procedures to extract the SN photometry. The host galaxywas not subtracted from the NIR frames; therefore the SN photometry contains an unknowngalaxy contamination component.The final CSP u ′ g ′ r ′ i ′ photometry of SN 2005gj is given in Table 3 and Y J HK s pho-tometry in Table 4. A minimum uncertainty of 0.015 mag in the optical bandpasses and0.02 mag in the NIR is assumed for a single measurement based on the typical scatter inthe transformation from instrumental to standard magnitudes of bright stars (Hamuy et al.2006).
3. Spectroscopy
The spectroscopic observations of SN 2005gj are summarized in Table 6. They wereobtained using five different telescopes/instruments at four observatories.A total of twelve spectra were obtained between early-October 2005 and late-January2006 with the Boller & Chivens CCD Spectrograph (CCDS) mounted at the MDM-2.4m tele-scope. This instrument uses a Loral 1200 ×
800 pixel CCD with 15 µ m pixel − and a 150 l/mmgrating (blazed at 4700 ˚A). We used a 2 ′′ slit which gives a dispersion of 3.1 ˚A pixel − in thewavelength range ∼ ′′ slit were used,reaching a dispersion of 3 ˚A pixel − in the wavelength range ∼ ′′ slit that gave a dispersion of 2.45 ˚A pixel − in the range ∼ × ′′ slit reaching a dispersion of 0.70 ˚A pixel − and 1.12 ˚A pixel − in the blue andred sides of the spectrum, respectively.Additionally we obtained two early spectra with the Double Imaging Spectrograph (DIS)mounted on the ARC-3.5m telescope at APO, and one spectrum with the Intermediatedispersion Spectrograph and Imaging System (ISIS) at the WHT-4.2m telescope at theRoque de Los Muchachos Observatory in La Palma, Spain. The DIS spectrograph has blueand red detectors, each uses a Marconi 2048 × µ m pixel − . Weused a 300 l/mm grating and a 1.5 ′′ slit which gives a dispersion of 2.4 ˚A pixel − . The ISISspectrograph has a blue and red arm, the blue arm using a EEV12 CCD and the red usinga MARCONI2 detector. We used a 300 l/mm grating in both the blue and red arm and a1 ′′ slit, which gives a dispersion of 0.86 ˚A pixel − in the blue and 1.47 ˚A pixel − in the red.Most of the spectra were obtained close to the parallactic angle to minimize relativechanges in the calibration of the blue and red parts of the spectrum due to differential refrac-tion through the atmosphere. The spectroscopic reductions were performed using standardIRAF tasks and included: bias and overscan subtraction, flat-fielding, combination of 2-4individual 2D spectra to reach the best signal-to-noise ratio in the final image, tracing andextraction of a 1D spectrum from the combined 2D image, subtracting the background skyaround the selected aperture, wavelength calibration using an arc-lamp, and flux calibration.In order to flux calibrate the spectra we observed 1-2 spectrophotometric standard stars pernight. The spectra from LCO, WHT and MDM were corrected by atmospheric telluric linesusing the spectrum of a hot spectrophotometric standard star and the spectra from APOwere corrected using a model atmosphere. This correction is not optimal for some of thespectra and there are evident residuals left in the corrected spectra.Figure 2 and 3 show a montage of the optical spectra of SN 2005gj obtained fromOctober 2005 to March 2006. We have split them in two figures to avoid crowding. Theposition of the most conspicuous spectral features have been indicated in this figure.
4. X-ray Observation
SN 2005gj was observed under Director’s Discretionary Time for 49.5 ks on 2005 Dec11/12 (ObsID 7241) with the
Chandra X-ray Observatory’s
Advanced CCD Imaging Spec- 10 –trometer (ACIS). The data were taken in timed-exposure mode with an integration time of3.2 s per frame, and the telescope aimpoint was on the back-side illuminated S3 chip. Thedata were telemetered to the ground in “very faint” mode.Data reduction was performed using the CIAO 3.3 software provided by the
Chandra
X-ray Center . The data were reprocessed using the CALDB 3.2.2 set of calibration files (gainmaps, quantum efficiency, quantum efficiency uniformity, effective area) including a new badpixel list made with the acis run hotpix tool. The reprocessing was done without includingthe pixel randomization that is added during standard processing. This omission slightlyimproves the point spread function. The data were filtered using the standard ASCA grades(0, 2, 3, 4, and 6) and excluding both bad pixels and software-flagged cosmic ray events.Intervals of strong background flaring were searched for, but none were found.Absolute
Chandra astrometry is typically good to ∼ . ′′
5, and we sought to improve itby registering the
Chandra image with an SDSS image.
Chandra point sources were foundusing the wavdetect tool, and their positions were refined using ACIS Extract version 3.101.Fourteen X-ray sources had SDSS counterparts, which we used to shift the
Chandra frameby a small amount (0 . ′′
15 in RA and 0 . ′′
07 in DEC). After the shift, the residual differencesbetween the
Chandra and SDSS sources had rms values of 0 . ′′
19 in RA and 0 . ′′
12 in DEC.We extracted counts in the 0.5–8 keV bandpass from the position of the supernova usinga standard extraction region ( ∼ ′′ radius), and we constructed response files with the CAIOtools and ACIS Extract. The background region is a source-free annulus centered on theposition of the supernova with inner and outer radii of 6 ′′ and 32 ′′ . Based on the 300 photonsdetected in this region, we expect 0.3 background counts in our source extraction region.We detect only two counts from the location of the supernova, but neither may beassociated with the supernova itself. The counts had energies of 4.0 keV and 6.5 keV, butone would expect some emission in the 0.5–2.5 keV range since this is where Chandra hasthe most collecting area. For example, the average effective area in each of the 0.5–2.5 keV,4.0–6.0 keV, and 6.0–8.0 keV bands is 470 cm , 270 cm , and 90 cm , respectively.We calculate our upper limits using the Bayesian method of Kraft et al. (1991). For the0.5–4 keV band, the 68% (95.5%) upper limit to the source counts is 1.14 (3.05). For the0.5–8 keV band, the 68% (95.5%) upper limit is 3.52 (6.14).Since no Type Ia supernova has been conclusively detected in X-ray, we have no a priori expectation of the spectral shape. We therefore adopt a simple absorbed power law with aphoton index of 2 and an absorbing column of n H = 7 . × cm − . For this choice of http://asc.harvard.edu
11 –spectrum, the count rate to flux conversion is 5 . × − erg/count in the 0.5–4 keV bandand 6 . × − erg/count in the 0.5–8 keV band.We therefore arrive at 68% (95.5%) upper limits on the X-ray luminosity of 1 . . × erg s − in the 0.5–4 keV band and 4 . . × erg s − in the 0.5–8 keV band. Basedon the above statements concerning the Chandra effective area, we feel the 0.5–4 keV limitis a more appropriate limit.
5. Results5.1. Optical light curves and colors
Figure 4 shows the early SDSS and MDM ugriz light curves combined with the late timecoverage given by CSP u ′ g ′ r ′ i ′ Y J HK s photometry. They give excellent multi-wavelengthoptical and NIR coverage and sampling of the first ∼
150 days after discovery ( ∼
140 rest-frame days after the time of explosion).In Table 5 we give important parameters derived from the light curves. We have a goodestimate of the time of explosion of the SN at JD 2,453,637.93 ± > σ in all filters at JD 2,453,639.94, and the last pre-discovery observation of the fieldat JD 2,453,635.91. The times and observed magnitudes at maximum in different filtersare presented in Table 5. They are calculated from a high order polynomial fit to eachlight curve. To estimate the errors we assume a gaussian distribution for the magnitudeuncertainty at each epoch and filter (assuming they are not correlated from epoch to epoch),then we draw randomly ∼ σ uncertainties are taken as the rms deviation of the mean valuescalculated from each simulated light curve.The risetimes, defined as the time between explosion and maximum light, become longerat redder wavelengths. They are 13.5, 19.7, 33.7, and 46.9 days in ugri filters, respectively,which correspond to 12.7, 18.5, 31.7, and 44.2 days in the rest-frame of the supernova.We also give in Table 5 the rest-frame magnitudes at maximum in different filters. Theyhave been corrected by Galactic extinction in the line of sight, using E ( B − V ) Gal = 0 . R V = 3 .
1, and K -corrections (see below), which are not negligible in this object since it is atredshift z ∼ . K -corrections have been calculated from the multi-epoch spectra presented in §
3. Inorder to estimate accurate K -corrections we need good spectrophotometric calibration in allthe wavelength range. Figure 5 shows the differences between the observed g − r photometriccolors obtained from the light curves and the synthetic colors calculated directly from thespectra, as a function of the observed photometric colors. We use an 8th order polynomialfit of the light curves to obtain the observed g − r color at the epoch of a given spectrum tobetter than ∼ .
03 mag. We can see that most of the spectra have good spectrophotometriccalibration in the wavelength range of g , r filters (3800–7000 ˚A), with a residual color of ∼ warp the spectra multiplying by a smooth function to match the observed col-ors, a technique commonly used for calculating K -corrections in Type Ia SN (Nugent et al.2002).First, we extrapolate the continuum in the blue and red sides of each spectrum presentedin § ugriz filters(2000–11000 ˚A). Using this extended version of the spectra we apply the CCM reddeninglaw iteratively until the synthetic g − r color matches the observed color in each spectrum.This procedure does not ensure that the calibration is good in the complete wavelengthrange; therefore we multiply by a smooth spline with knots at the effective wavelength ofthe SDSS filters until the synthetic u − g , g − r , r − i and i − z colors match the observedcolors obtained, again by using polynomial interpolation of the light curves.The K -corrections for the same filter (Hamuy et al. 1993) calculated from the modified,spectrophotometrically calibrated spectra using SDSS passbands are listed in Table 7. The K -corrections are probably accurate to ± g and r filters, and to ± u , i and z where we had to extrapolate the spectra. However, our estimate is not preciseand we can not exclude the possibility of even larger errors. After fitting with a low orderpolynomial we use the results to transform the observed SDSS magnitudes in Table 2 to therest-frame.We corrected the CSP u ′ g ′ r ′ i ′ magnitudes to rest-frame SDSS ugri magnitudes usingcross-filter K -corrections according to the prescription of Nugent et al. (2002). These takeinto account the difference between the CSP and SDSS passbands convolved by the SED ofSN 2005gj and allow us to put all the rest-frame magnitudes in the same system. We findthat the differences between the same- and cross-filter K -corrections are small ( . .
03 mag)at all epochs.Figure 6 shows the evolution of the rest-frame colors of SN 2005gj as a function of time 13 –after explosion. We have corrected the magnitudes by K -corrections and Galactic extinctionin the line of sight. As a comparison we also plot the color evolution of the overluminousType Ia SN 1991T, a typical Type IIn SN 1999el (Di Carlo et al. 2002) (both obtainedfrom spectral templates of P. Nugent ), and two previous cases that are thought to beType Ia explosions in a very dense environment: SN 2002ic and SN 1997cy. For SN 2002icwe use the published BV I photometry (Hamuy et al. 2003) corrected by Galactic extinction( E ( B − V ) = 0 .
06; SFD). We calculate cross-filter K -corrections using the calibrated spectraof SN 2005gj to transform observed magnitudes of the SN at z = 0 . B → u , V → g , I → i . The time of explosion of SN 2002icis assumed to at ≈ JD 2,452,581.5 (2002, Nov. 3 UT; Deng et al. 2004). We obtain rest-frame colors in the SDSS system for SN 1997cy (Germany et al. 2000; Turatto et al. 2000) bytransforming the K -corrected V RI magnitudes in Germany et al. (2000) to gri magnitudesusing S -corrections calculated directly from spectra of SN 1997cy available online in theSUSPECT database . We supplement this with synthetic colors from the spectra. We alsocorrect by Galactic extinction in the line of sight ( E ( B − V ) = 0 .
02; SFD) and assume thetime of explosion of SN 1997cy to be JD 2,450,582.5 (1997, May. 14 UT; Germany et al.2000), which is very uncertain, since it is taken as the time of detection of the gamma-ray burst GRB 970514, which may not have been associated with the SN. Magnitudes forSN 2002ic and SN 1997cy are not corrected by extinction in their host galaxies, which isunknown.Initially the evolution in rest-frame u − g and g − r colors of SN 2005gj, up to ∼ ∼ ∼ . g − r , and evolves to redder colors at latertimes. SN 1991T reaches its maximum colors of ( u − g )=1.5 mag and ( g − r )=1.0 mag at ∼
50 days (30 days after maximum), and after that it enters the nebular phase and becomesbluer. At late times the u − g color of SN 2005gj has a slow linear increase and becomessystematically redder than SN 1991T at >
70 days, while the g − r color stays approximatelyflat at ( g − r )=0.5 mag (and bluer than SN 1991T) between 60 −
110 days. The evolutionin u − g color of SN 2002ic is very similar and consistent with SN 2005gj. SN 1997cy has asimilar color evolution but it is ∼ g − r color between 60 −
100 days. However,Germany et al. (2000) give K -correction errors of ∼ S -corrections, therefore this does not imply a significant difference between SN 1997cy andthe other two Type Ia/IIn supernovae. http://supernova.lbl.gov/ ∼ nugent/nugent templates.html http://suspect.nhn.ou.edu/ ∼ suspect/
14 –The rest-frame r − i and i − z color evolution of SN 2005gj is very different fromSN 1991T and closely follows the evolution of a SN IIn. The r − i colors of SN 1997cy arealso consistent with SN 2005gj. From these comparisons it is clear that the colors, a proxyfor the temperature of the photosphere, of SN 2005gj and two earlier cases of SN Ia stronglyinteracting with their circumstellar medium are dominated by the radiation coming from theejecta-CSM interaction.In Figure 7 we present the ugri light curves of SN 2005gj in absolute magnitudes. Wealso show the light curves of SN 1991T, SN 2002ic and SN 1997cy obtained from the literatureand corrected to SDSS rest-frame magnitudes as explained above.SN 2005gj has peak ugriz absolute magnitudes in the range − . − . ∼ . − . u light curve is consis-tent with a linear decline after peak luminosity at a constant rate of 0 . ± .
001 mag day − .The g light curve has a ∼
20 day plateau with roughly constant luminosity after maximum(20 −
40 days after explosion), then the light curve declines linearly between 40 −
100 days at0 . ± .
001 mag day − and continues its linear decay at later times, but with a shallowerslope of 0 . ± .
002 mag day − . The r and i light curves have a similar plateau shape be-tween 20 −
60 days and a constant linear decay at later times of 0 . ± .
001 mag day − . Thechange in slope observed in the g -band at late times is less clear in ri , but still present. Thesecondary maximum present in ri light curves of SN 1991T and other SN Ia is completelyabsent in SN 2005gj.The light curves of SN 2002ic are fainter than SN 2005gj at all times by 0 . − . . − . u and g light curves of SN 2002ic are intermediate betweenthose of SN 1991T and SN 2005gj until around 40 days after explosion; after day 50 whenthe ejecta-CSM interaction had become dominant in SN 2002ic (Hamuy et al. 2003), theyclosely resemble the decline rates of SN 2005gj in the same bands. The i band light curveof SN 2002ic showed definite evidence for a weak secondary maximum, which is again inter-mediate in morphology between the strong secondary maximum observed in SN 1991T, andthe absence of such a feature in SN 2005gj. The gr light curves of SN 1997cy are consis-tent with linear decay of ∼ .
008 mag day − , being ∼ . ∼
60 days after explosionand there is no information near peak to compare with the luminosities of SN 2005gj andSN 2002ic. However, if we extrapolate the gr light curves to the time of peak luminosity wefind that the absolute magnitudes at maximum of SN 1997cy would be within ∼ . We used the CSP
J HK s photometry obtained between 59-166 days after the explosion(rest-frame) to construct the absolute magnitude light curves of SN 2005gj. The observedmagnitudes were corrected by Galactic reddening in the line of sight ( A J = 0 . A H =0 . A Ks = 0 .
044 mag) and K -corrections. We calculated K -corrections for the samefilter using the spectral templates of P. Nugent for the Type IIn SN 1999el (which are derivedfrom black-body curve fits to the photometry), because as we showed in § r − i and i − z ). The values of the K -correctionsare consistent with being constant in this time range: K J ≈ − .
12 mag , K H ≈ − .
14 mag, K K s ≈ .
16 mag.The
J HK s absolute magnitude light curves of SN 2005gj are presented in Figure 8.For comparison we show the NIR light curves of a normal Type Ia obtained from syntheticphotometry of spectral templates and the Type IIn SN 1999el (Di Carlo et al. 2002). TheType Ia light curves have been shifted in magnitudes to match the mean absolute magnitudeat maximum of SN Ia (Krisciunas et al. 2004). SN 2005gj is 1 . − >
60 daysafter explosion), with decline rates of ∼ .
014 mag day − in J , ∼ .
013 mag day − in H ,and ∼ .
011 mag day − in K s . These values are similar to the decline rates in the optical ri bands.Since there are no template images of the host galaxy obtained in the NIR bands, thelight curves are preliminary and the analysis has to be taken with caution. The SDSS and CSP magnitudes were used to produce a quasi-bolometric light curveof SN 2005gj covering the optical wavelength range from 3000 − ,
000 ˚A( u → z ). Wecorrected the magnitudes by Galactic extinction in the line of sight and K -corrections toobtain magnitudes in the rest-frame ugriz filters (see § AB system obtained from the SDSS website . The AB magnitudes derived in this wayare transformed directly to bandpass averaged fluxes using the definition of the AB system .
16 –(Oke & Gunn 1983) and they are assigned to the frequencies that correspond to the effectivewavelengths of the SDSS ugriz filters, calculated from the filters using the definition inFukugita et al. (1996): 3567, 4735, 6195, 7510, and 8977 ˚A. We use the trapezoidal ruleto obtain the integrated flux from u to z , this is between λ = 3340 ˚A and λ = 9596 ˚A,where the limits of the wavelength coverage are obtained from λ = λ eff,u − ∆ λ u / λ = λ eff,z + ∆ λ z /
2. We extrapolate linearly the u and z light curves at late times to fillin the lack of coverage of the SDSS in z , and CSP in u and z bands, including the MDM z -band data at these epochs. The integrated fluxes are converted to luminosity assuminga luminosity distance to SN 2005gj of 268.5 Mpc and a spherically symmetric distributionof the output energy. We present the integrated, quasi-bolometric luminosities from u to i ( L ( u → i ) ) and from u to z ( L ( u → z ) ) as a function of time in Table 8.In order to estimate bolometric UVOIR luminosities we calculate time dependent bolo-metric corrections to include the energy output of the SN at wavelengths bluer than u -band( λ < z -band ( λ > χ -minimization to fit the optical SED with a two parameterblack-body function: temperature and a multiplicative scaling factor. The scaling factor is acombination of fundamental constants and the square of the angular radius of the sphericalblack-body: ( R bb /d ) , where R bb is the radius of the black-body and d is the distance tothe SN. We calculate time-dependent bolometric corrections by integrating the black-bodydistributions in the ultraviolet and NIR/IR regions, and converting the integrated fluxes toluminosities as explained above. For the CSP data we also include NIR flux densities derivedfrom the reddening and K -corrected Y J HK s magnitudes after fitting the light curves withlow order polynomials.At early phases before peak the bolometric corrections account for ∼ −
65% of thetotal integrated luminosity, of which 85 −
93% is from the ultraviolet part of the spectrumand only 7 −
15% from the NIR/IR. As the supernova evolves, the ejecta expands andshocks the circumstellar gas. The energy emitted in the ultraviolet/blue part of the spec-trum declines quickly after maximum light and most of the energy is emitted in the opticalregion, coinciding with the appearance in the spectrum of emission/absorption features ofthe intermediate and iron-peak elements (see § −
150 days after explosionthe bolometric corrections account for ∼ −
45% of the total output luminosity, with theNIR/IR correction dominating completely over the blue/ultraviolet at >
60 days.The bolometric UVOIR luminosities, black-body temperatures and radii derived fromthe fits are presented in Table 8. The uncertainties in black-body temperatures and radii arecalculated using the diagonal terms of the covariance matrix obtained from the χ minimiza- 17 –tion. We add a 10% error in the distance to the SN to the error in the black-body radii, whichcomes from the random and systematic uncertainties in the value of the Hubble constant(Freedman et al. 2001; Riess et al. 2005). The uncertainties in the bolometric luminositieswere estimated by propagating errors through the trapezoidal integration of the SED, takinginto account: uncertainties in the photometry, light curve interpolation and fitting, Galacticextinction, K -corrections, and distance to the SN. To approximately take into account theerrors introduced by the bolometric corrections we multiplied these values by p χ ν when thereduced χ is greater than 1.In Figure 9 we show some examples of black-body fits to the optical SED at differentepochs. At early times, shortly after explosion, the SED is very well fit by a hot ∼ ∼ χ ν of the fits (see Table 8).Figure 10 shows the bolometric light curve of SN 2005gj in the top panel and the evolu-tion in temperature and radius from the black-body fits in the lower panels. The early dataof the bolometric light curve are well fit by an exponential rise in luminosity, L ( t ) ∝ e . t .The time of maximum bolometric luminosity occurs between 6.6–18.8 days after explosion.After maximum, the bolometric light curve is very well approximated by an exponentialdecay in luminosity, linear in the logarithmic scale shown in Figure 10, L ( t ) ∝ e − . t (0 .
014 mag day − ). This is consistent with the exponential density distribution of the ejectaof Type Ia SN (Dwarkadas & Chevalier 1998), whereas the distribution of ejecta around core-collapse supernovae is better approximated by a power-law (Chevalier & Fransson 2003).Extrapolating the pre- and post-maximum fits we get a maximum bolometric luminosity of L maxbol = 5 . × ergs − , which is ∼ ∼ In Figure 11 we show a comparison of the spectra of SN 2005gj with spectra of SN 2002icand SN 1997cy obtained at similar times after explosion. The spectra of SN 2005gj andSN 2002ic are very similar at all times. They are characterized by strong and broad 18 –Hydrogen-Balmer lines H α and H β in emission and a blue continuum at early times thatbecomes redder and increasingly dominated by absorption/emission P-cygni profiles fromFe-peak ions (e.g., Fe II , Fe III , Ni
III , Si II , S II ). Benetti et al. (2006) proposed that SN 2002ic-like events are well explained by the core-collapse of a massive star in a dense medium, casting doubt in the previous classificationof SN 2002ic as a Type Ia supernova. The authors found relatively good agreement at alltimes between the spectra of SN 2004aw (Taubenberger et al. 2006), a Type Ic supernova,and SN 2002ic.We used the SuperNova IDentification code, SNID (Matheson et al. 2005; Miknaitis et al.2007; Blondin & Tonry 2007), to find the spectra that best match SN 2005gj at differentepochs. SNID cross-correlates an input spectrum with a library of supernovae spectra. Inthe library we included spectra of 5 normal SN Ia, two 1991T-like objects, two 1991bg-likeobjects, 4 broad-lined SN Ic (or hypernovae), and 3 normal SN Ic (including SN 2004aw),that were chosen to span a wide range of observed properties of SN Ia and SN Ic. In Table 9we present the supernovae and the epochs of the spectra in the library. We fixed the redshiftof SN 2005gj at z = 0 . z = 0 . As shown in Figure 2, H γ is also visible in the early spectra of SN 2005gj.
19 –SN 2005gj, and a fourth order polynomial. A normal SN Ia does not fit as well as SN 1991T.This procedure is very similar to the fits to SN 2002ic (Hamuy et al. 2003) and SN 2005gj(Aldering et al. 2006) presented in previous studies. In Figure 13 we show examples of thespectra decomposition at four epochs. We excluded from the fit a region of ±
100 ˚A aroundthe H α and H β lines and obtained a good fit for the remainder of the spectrum. We analyzed the Balmer emission features in the spectra using the sum of two Gaussiancomponents to model the line profiles. This decomposition gives much better fits for H α at all epochs than a single Gaussian and it is physically motivated (Chugai 1997a,b). Thespectra of Type IIn SN show Balmer features with both a narrow and broad component thatcan be explained by radiation coming from different regions of the ejecta/CSM, whether it isdirect emission from the shock-heated CSM (broad component) or emission from un-shockedgas photoionized by the SN radiation (narrow component). The H β line is unresolved oronly marginally resolved for most of the spectra. Therefore a single Gaussian component wasused to fit the line profile. We used a third order polynomial to model the local continuumaround each line that was included in the Gaussian fits. It is important to stress that at latetimes there is a broad Fe II feature intrinsic to the supernova spectrum in the region of H α (see spectra in Figure 13) that makes the definition of the continuum less reliable and mayaffect the line measurements.The results of the Gaussian fits to the H α and H β emission features, integrated fluxes andFWHM, are shown in Table 10 as a function of epoch of the spectra. We have excluded thetwo spectra with better resolution because they show P-Cygni profiles (see below). We usedthe flux calibrated spectra corrected to match the observed g , r magnitudes (as explainedin § α and H β luminosities (top right and bottom left panels) and the Balmer decrement (bottom rightpanel). The FWHM of H α varies between ∼ −
500 km s − (narrow component) and ∼ − − (broad component), with the broad component showing a slow increasein time. The FWHM of H β varies between ∼ − − and does not show evidentevolution. The luminosities of H α -narrow and H β lines evolve in a similar fashion, increasing 20 –at early times to peak at ∼
12 days with luminosities 6 − . × ergs − , then they decayand stay roughly constant after 50 days. The evolution of H α -broad is similar during the first50 days, peaking at 1 . × ergs − , but it shows an increase at later times. Compared withthe H α luminosities observed in SN 2002ic, both components are ∼ α (sum of narrow and broad components)and H β fluxes, stays approximately constant during the first 30 −
40 days (mean = 2 . .
5) and is consistent with the theoretical value in Case B recombination of H α /H β =2 .
86 (Osterbrock 1989). At later times it shows an steady increase, reaching H α/ H β ∼ − ∼
80 days. In Case B recombination a Balmer decrement H α/ H β > .
86 is usuallyinterpreted as evidence for the presence of internal extinction in the host; however, thelarge values observed at late times would produce an Na I D interstellar absorption doubleteasily detectable in the spectra, which is not observed (see also Aldering et al. (2006)), andthe evolution in time is not expected. Aldering et al. (2006) proposed that the H levelpopulations are in Case C recombination, where the optical depth in the H α line is highimplying high densities and greater importance of collisional processes. In this scenario, theobserved change in the Balmer decrement could indicate that collisional excitation becomesincreasingly important at later times (Branch et al. 1981; Turatto, et al. 1993). SN 2002ic(Deng et al. 2004) and other SN IIn, like SN 1988Z (Aretxaga et al. 1999) and SN 1995G(Pastorello et al. 2002), have also shown large values of the Balmer decrement and thereforemay have similar physical processes affecting the formation of the Balmer lines.In Figure 15 we show the regions around H α and H β features in the best resolutionspectra from ISIS and LDSS-3, obtained at 44 and 115 days after explosion, respectively. Weclearly detect P-Cygni profiles in all these features, which indicates the presence of an outflowmoving at ∼ − ; however, these measurements are limited by the resolution of thespectra between ∼ − (FWHM). After correcting for the resolution we obtain anoutflow velocity of 60-70 km s − . The detection of P-Cygni-like absorption rules out an H II region in the line of sight that could be producing the narrow emission/absorption features;the line profiles are intrinsic to the SN. Aldering et al. (2006) detected P-Cygni profiles inHe I , H α and H β , in a high resolution spectrum obtained with LRIS+Keck 71 days afterthe explosion. They derived a wind velocity of v w ≈
60 km s − consistent with our estimate.Kotak et al. (2004) also detected a P-Cygni profile in the a spectrum of SN 2002ic obtained256 days after explosion. 21 – We used the parameterized resonance-scattering code SYNOW (Fisher et al. 1997; Fisher2000) to identify the lines in the spectra obtained near maximum light of SN 2005gj. SYNOWis a fast supernova spectrum-synthesis code used for direct (empirical) analysis of supernovaspectra, mainly to identify the lines, their formation velocities and optical depths. Thecode is based on simple assumptions: spherical symmetry, homologous expansion, a sharpphotosphere that emits a black-body continuous spectrum, and line formation by resonance-scattering, treated in the Sobolev approximation. We have used the latest version of thecode that includes a Gaussian distribution of optical depths.Figure 16 shows the spectrum of SN 2005gj at 17 days after explosion (2 days before g maximum) and the best synthetic spectrum obtained with SYNOW. We also show forcomparison the spectrum of SN 1991T obtained at -3 days with respect to the time of B maximum. The spectra have been locally normalized as in Jeffery et al. (2006). Thesynthetic spectrum has a black-body continuum temperature T bb = 11000 K, photosphericvelocity v phot = 10000 km s − , and excitation temperature T exc = 10000 K. We find areasonably good match with the spectrum of SN 2005gj using the following lines/multiplets:Fe III λ λ III λ III , S II λ λ II λ III features and weak S II doublet and Si II (Jeffery et al. 1992; Mazzali et al. 1995; Fisher et al. 1999).The main discrepancy between the SYNOW modeling of SN 2005gj and SN 1991T isin the optical depths of the lines. The fit to SN 2005gj needs unphysically small opticaldepths, approximately 1/10th of the values used for SN 1991T around maximum light. Weinterpret this as an effect of the extra continuum radiation that is added by the ejecta-CSMinteraction, which is veiling (Branch et al. 2000) the supernova lines (e.g., Hamuy et al. 2003;Aldering et al. 2006). This interpretation is supported by the good agreement obtained fromfitting the spectra of SN 2005gj using a simple polynomial continuum added to the spectraof SN 1991T at the same epochs after explosion (see Figure 13).
6. Discussion
We have presented extensive spectroscopy and optical/NIR photometry of SN 2005gj ob-tained by the SDSS-II and CSP supernova groups during the first ∼
150 days after explosion,and also an X-ray observation at 74 days that gives an upper limit on the X-ray luminosity.We have shown the remarkable similarity in spectroscopic and photometric properties be- 22 –tween SN 2005gj and SN 2002ic, which is thought to be the first clear case of a thermonuclearsupernova explosion embedded in a dense CSM. The observational properties of SN 2005gjsupport this interpretation, they are summarized as follows: • Spectroscopic evidence for a shock propagating into an Hydrogen-rich medium closeto the site of the explosion inferred from the presence of Balmer lines with narrow(FWHM ∼ −
500 km s − ) and broad (FWHM ∼ − − ) componentsat all times. The Balmer lines show P-Cygni profiles in the highest resolution spectraobtained at 44 and 115 days after explosion, these detections show the presence of aslow ( ∼
100 km s − ) moving outflow. Both observations support the interpretation ofthe supernova ejecta interacting with a dense circumstellar material. • Spectrum evolves from a very blue continuum (13000 K black-body) similar to SN IInat ∼ III , weak S II andSi II ) and the spectra are well matched by the overluminous Type Ia SN 1991T diluted with a polynomial continuum at similar times after explosion. • Very luminous and slowly declining bolometric light curve. The linear decay in lumi-nosity after peak ( ∼ .
014 mag day − ) suggests an exponential density distribution ofthe ejecta, which is consistent with the ejecta-density profiles obtained from simulationsof SN Ia.The data presented here on SN 2005gj makes the interpretation of 2002ic-like eventsas thermonuclear supernovae in a dense CSM, initially proposed by Hamuy et al. (2003),stronger. In contrast with Benetti et al. (2006), we find that the overall shape of the spectraof SN 2005gj are more consistent with spectra of SN Ia at different epochs. Specifically,Type Ic SNe usually do not show the S II doublet at ∼ II near maximum(Taubenberger et al. 2006). This is one of the identifying features in SN Ia spectra, alsopresent in the overluminous SN 1991T (Phillips et al. 1992). In the spectrum of SN 2005gjobtained at 17 days (see Figure 16) we detect a weak double absorption that we identify withS II , that is much stronger in the spectrum of SN 2002ic around maximum light. We cansee on the top of Figure 11 that the spectrum of SN 2002ic obtained 24 days after explosionclearly shows this feature. Other conspicuous features observed in SN 2005gj and SN 2002icaround maximum are Fe III and Si II . These features are present in SN 1991T, but Fe III isnot observed and Si II is generally weaker in SN Ic. 23 –SN 2005gj has stronger ejecta-CSM interaction than SN 2002ic. The peak bolometricluminosity is ∼ . α are ∼ diluted by the stronger continuum. The absence of evidence fora secondary maximum in SN 2005gj, whereas the i band light curve of SN 2002ic does showa hint of such a feature, is likewise consistent with the ejecta-CSM interaction in SN 2005gjhaving been stronger than in SN 2002ic. The circumstellar interaction of core-collapse supernovae with a circumstellar mediumhas been studied in detail in the literature (see Chevalier & Fransson (2003) for a review).When the fast moving ejecta encounters the approximately stationary CSM, a forward shockmoving into the CSM (also called circumstellar shock) and a reverse shock develops. Thefast-moving shockwave implies large post-shock temperatures, therefore radiating energy inthe X-ray regime. The density distribution of the ejecta and the CSM can be well describedby power-laws in radius, which leads to a set of self-similar analytical solutions for theevolution of the shock radius in time (Chevalier 1982). The physics of the ejecta-CSMinteraction in the case of thermonuclear supernovae is basically the same, the main differenceis in the distribution of the ejected material which follows an exponential function in velocity(Dwarkadas & Chevalier 1998). In this case the solutions are no longer analytic. The densityprofile of the shocked region is different in the case of exponential ejecta expanding into aconstant density medium, but the similarity increases for expansion into a wind profile whosedensity decreases as ∝ r − .A simple self-similar model of a SN shock expanding into a medium with a power-law density decline, as suggested for core-collapse SNe by Chevalier (1982), is ruled outfor this object by several observations: the exponential decrease in luminosity, suggestingan exponential ejecta density profile; the strange behavior of the broad and narrow H α components; and the decrease in the blackbody radius at later times. Detailed calculationsof the SN-CSM interaction would require highly detailed hydrodynamic modeling, whichare beyond the scope of this paper. Instead herein we focus on trying to explain the basicfeatures of SN-CSM interaction as deduced from the observational data.The initial velocity of a SN shock wave as it breaks out from the surface is at least of theorder of 2 × km s − . The broad H α velocities that are seen in the first week or so are ofthe order of 1500 km s − , and increase to more than twice this value after ∼
50 days. These 24 –velocities are almost an order of magnitude smaller than expected SN blast wave velocitiesin the early stages, and a factor of few smaller even after ∼
50 days. Furthermore, theSN shock velocity would be expected to gradually decrease as the shock moves outwards,whereas the H α profile actually indicates an increasing velocity after ∼
50 days.For these reasons, we suggest that the broad H α lines do not indicate the SN velocity.Instead, we put forward a scenario of a shock expanding into a two-component ambientmedium: a low density wind in which are embedded high-density clumps. In this picture,there should theoretically exist three different velocity components: a broad velocity com-ponent, which is not easily seen in this case, and is related to the velocity of the blast waveitself; an intermediate velocity component, which is what we have referred to as the broadH α and is related to the velocity of the shock driven into the clumped material; and anarrow velocity component, which may be related to the narrow H α and is representativeof the velocity of the ambient medium. This scenario is like the scenario put forward by(Chugai & Danziger 1994) to explain the origin of the broad, intermediate and narrow linecomponents in SN 1988Z. The large H α luminosity of SN 2005gj at late times is very sim-ilar to that seen in other Type IIn SNe, and is especially large considering that this wasa Type Ia. However, there are significant differences. We do not see a really broad linecomponent representative of the SN velocity, although there are some suggestions that thismay be appearing at late times. In particular, the H α profile of the spectrum obtained at ∼
150 days is better fitted by three components, including a very broad component withFWHM ≈ − .Our scenario envisions the Type Ia SN shock wave expanding in a clumped mediumpresumably formed by mass-loss from a companion star. The broad component is not easilyvisible in H α initially because the forward shock is not radiative. The density of the clumpsis much higher than that of the interclump (ambient) medium. When the SN shock waveinteracts with a dense cloud or clump, it drives a strong shock into the clump. A reflectedshock is driven back into the expanding ejecta (Klein et al. 1994). Assuming pressure equi-librium, the ratio of the velocity of the clump shock to that of the blast wave is inverselyproportional to the square-root of the ratio of the clump density to that of the interclumpmedium. The optical emission arises from behind the clump shock, probably by reprocessingof the X-ray emission.In this model, the intermediate component represents the velocity of the clump shock,which is probably radiative. If we assume that the initial velocity of the SN shock waveis ∼ − and the broad H α emission velocity is ∼ − , then the ratio ofvelocities is 13 −
14. This indicates that the clump density is about 14 , or ∼
200 times theinterclump density. Note that the optical emission, which goes as density squared, will then 25 –be 200 times, or about 40,000 times greater compared to that from the interclump medium.This is consistent with the fact that no broad line emission is seen from the interclumpmedium. If the initial velocity is much higher, as is conceivable, the clump density couldbe up to ∼
50% higher, and the ratio between the emission from the dense clumps andinterclump medium even larger.What value of the clump density is suggested? A shock wave traveling at 1500 km s − would be radiative if it were expanding in a medium whose density is greater than ∼ cm − , whereas a 2500 km s − shock would require minimum densities of the order of10 cm − (Draine & McKee 1993) in order to be radiative. The CSM density, being twoorders of magnitude smaller, would then to be & cm − . These are just minimum values,and it is conceivable that the actual clump density is much higher. This result is consistentwith the conclusion of Aldering et al. (2006).The observations show that the broad H α width increases after 50 days, suggestingan increase in the clump shock velocity at later times, which could perhaps be due to adecrease in the clump density. Conversely, however, the luminosity of the H α also increases,suggesting an increase in the electron density. At the same time, we would expect the SNshock to be decreasing in velocity as it continues its outwards expansion.We suggest that the way to reconcile these observations is a scenario in which the densitywithin the clump medium starts out higher than 10 cm − , probably as high as 10 cm − inthe first few days, and decreases gradually outwards. The almost constant behavior of theFWHM of the broad H α suggests that the density profile of the ambient medium is flatterthan r − . Since we want the clump shock to be radiative even when the shock velocityis almost 3000 km s − , this suggest that the density at ∼
150 days is greater than about10 cm − . And since the density is decreasing outwards, we infer that the density close in iseven larger. Over the entire period of observations the clump density is large enough thatthe shock driven into these clumps is always radiative. The density of the ambient mediumis two orders of magnitude smaller, as discussed above. The high bolometric luminosity isconsistent with these values.For the first ∼
50 days the H α emission arises only from the radiative shock driven intothe dense clumps. However, by ∼
50 days the SN forward shock, which is decreasing invelocity, enters the radiative regime, and the cooling shell of material begins to contributeto the H α luminosity. The velocity of the SN shock is quite large, and its contributioninitially is not a large fraction of the total H α luminosity. But as it expands outwards,its velocity decreases and the shock becomes more radiative, and the contribution to thetotal H α luminosity increases, more than compensating for the decreasing density. If thisconclusion is correct, then we would expect that a broad velocity component would be visible 26 –in the H α spectra, whose intensity would gradually increase with time even as the FWHMdecreases. Although the underlying supernova contamination makes it hard to isolate abroad component, it is suggestive that by day ∼
150 the spectrum is best fit by a third,much broader component of the velocity, thus providing support for this line of reasoning.Finally, in this model the narrow line emission arises from the unshocked slowly expand-ing ambient material, presumably the outflow that we find expanding at ∼
60 km s − . Wenote that although the width of the narrow line H α emission as listed is higher, it is stillunresolved, and it is possible that within the limits of resolution the narrow line componentand outflow velocity are indeed the same.To summarize, in this model the Type Ia SN expands in a clumped ambient medium,with the clump density about ∼
200 times that of the surrounding medium close in tothe star, and decreasing as we go outwards. The H α emission initially arises mainly fromthe shock driven into the dense clumps. The SN shock propagating into the interclumpmedium begins to enter the radiative regime around day 50, and its contribution to the H α emission gradually increases beyond that coming from the clumped medium, leading to thegradual rise in the H α emission. We note that several features of this model are similarto the model presented by Chugai et al. (2004) for SN 2002ic, thus further supporting thesimilarity between the two supernovae.The upper limit on the X-ray luminosity obtained at 74 days after the explosion canput a constrain on the mass loss rate from the precursor or companion (e.g., Immler et al.2006). Assuming that the X-ray luminosity is dominated by emission from the reverse shockwe obtain ˙ M . × − M ⊙ yr − (2 σ ) using Equation 3.10 in Fransson et al. (1996). Thisvalue has to be taken as an approximate estimate because we are making several assumptionsabout the physical properties of the ejecta-CSM interaction that should be calculated usingdetailed hydrodynamical simulations: a constant velocity of the shock, V sh ≈ − ;a solar composition of the CSM material; an electron temperature at the reverse shock of T e = 10 K, which comes from the modeling of SN 2002ic (Nomoto et al. 2005); a flatdensity profile of the CSM, ρ ∝ r − ; and a power-law ejecta density profile with index n = 7(Nomoto et al. 1984).We can also estimate a mass loss rate from the companion using the density of theambient medium ( n ∼ cm − ), the initial optical radius of the CSM ( R ≈ cm), andthe velocity of the wind: ˙ M = 4 π R v ρ , this is assuming a flat density profile for the CSM.We obtain: ˙ M ≈ × − M ⊙ yr − , which is in agreement with the 2 σ upper limit calculatedfrom the X-ray luminosity.The presence of Balmer lines in emission in the first spectrum obtained 6.6 days af- 27 –ter explosion shows that the ejecta started to interact with the CSM at an earlier epoch(Aldering et al. 2006). Extrapolating linearly to zero flux the early increase of H α and H β fluxes we find that the ejecta-CSM interaction started 3 ± R i ≈ . × cm. The outer radius of the CSMcan be estimated assuming a constant velocity of the shock of V sh ≈ − over thefirst year. We detect H α in emission in a spectrum obtained at 368 days after explosion,which will be presented elsewhere, putting a lower limit on the outer radius of the CSM, R o & × cm. This is also consistent with a Type Ia SN with an exponential ejectadensity profile expanding outwards in a medium of average density & cm − .In the interpretation above we assume that the broad component of the HydrogenBalmer lines originate in the dense clumps, while the narrow component arises from thephotoionized un-shocked gas. However, Thompson scattering of the lines has been consideredas an alternative mechanism that can explain relatively well the symmetric line profiles ofSN 2002ic (Wang et al. 2004) and SN 2005gj (Aldering et al. 2006). In this scenario, bothcomponents would arise from a single high density region. The total mass of hydrogen inthe emitting region would be M H ≈ × − (10 /n e ) M ⊙ , where n e is the average electrondensity in the emitting zone, as calculated from the luminosity of the H α line at maximumusing the Case B recombination coefficient. The electron density must be sufficiently high, n e ≈ cm − , to be consistent with the line ratios of He lines observed in the spectra(Aldering et al. 2006), and a high electron density would explain the non-detection in X-rayand radio (Soderberg & Frail 2005). However, it is unlikely that this model would be ableto explain the initial constancy and then rise of the broad H α luminosity. The SDSS-II Supernova Survey has a well understood discovery efficiency of SN Ia atlow redshift ( z . . z < .
12, one photometric identification, and the spectroscopically confirmed peculiar events:SN 2005hk (Phillips et al. 2007) and SN 2005gj. Since the detection efficiency of 2002ic-likeobjects has not been carefully modeled, we can only put a lower limit on the fraction of theseevents. The spectroscopic confirmation of one object at z < .
12 puts a lower limit of 5 +7 − %(68% confidence) in the fraction of 2002ic-like events among SN Ia at low redshift. From thepreviously known (2002ic) and probable events (1997cy and 1999E) the estimated fraction 28 –is ∼
1% of SN Ia discovered between 1997 and 2002, which is consistent with our limit.In the fall of 2006 we obtained the spectrum of a slowly declining supernova that wasdiscovered in 2005, but did not have a spectroscopic classification, SN 7017 . To our surprise,the late spectrum of SN 7017 resembles that of SN 2005gj one year after explosion and theearly photometry also shows similarities which lead us to classify it as the highest redshiftSN 2002ic-like object observed to date, at z = 0 .
27 (Prieto et al. 2007, in preparation).Considering SN 7017 in the spectroscopically confirmed sample of SN Ia during the 2005season, a total of 129 SNe at z . .
42, we have that 2/130 (1.5%) are SN 2002ic like objects,which is consistent with the low limit on the fraction at low redshift estimated before.However, this fraction has to be taken with extreme care and probably does not reflectthe true fraction. This is because the discovery efficiency of SN Ia declines as a functionof redshift and the total number of spectroscopically confirmed SN Ia does not include SNewith good Ia-like light curves that were not spectroscopically classified. A more careful studyof the rates of SN 2005gj-like supernovae in the SDSS-II is planned for a future publication.The host galaxies of supernovae can provide important clues about their progenitors.The host of SN 2005gj is a very blue, low-luminosity dwarf ( M B ≈ − Z < . Z ⊙ , with a burst of star formation ∼
200 Myr ago. SN 2002ic has alate type (Sbc) spiral host with a well defined core. The host of SN 1997cy is also a blue, low-luminosity ( M V ≈ − . ∼ − . L ∗ galaxiesobserved by GALEX at redshift z < . L ∗ galaxy inthe K -band (Kochanek et al. 2001). SN 7017 at redshift z = 0 .
27, has a blue, dwarf-likehost galaxy with absolute magnitude in B of − . This is the internal name given by the SDSS-II Collaboration. It was not announced in an IAU circularbecause of the late spectroscopic classification.
29 –which indicates they are low metallicity systems. For example, a dwarf galaxy with intrinsicluminosity M B = −
18 has an Oxygen abundance of 12 + log(O / H) ≈ . K -band luminosity, that whenconverted to metallicity using the luminosity-metallicity relationship derived by Salzer et al.(2005), makes it consistent with the solar value. The host luminosities are only an approx-imate indicator of their metallicities, therefore spectra of the hosts are needed to infer themetallicities and star formation rates (SFRs) of these galaxies. However, it is interestingto note that the range of host galaxy properties of SN 2002ic-like events seem to be incon-sistent with the host galaxies of GRBs associated with supernovae (Stanek et al. 2006) andbroad-lined type Ib/c SNe (Modjaz et al. 2007).Type Ia supernovae are observed in all types of galaxies. There is a well establishedcorrelation between the morphology of their host galaxies and the peak luminosity of theSNe: brighter supernovae (1991T-like) tend to explode in late type spirals and irregulars withrecent star formation, while intrinsically fainter events (1991bg-like) are observed mainly inearly type galaxies with an old stellar population (Hamuy et al. 1995, 1996a; Branch et al.1996; Hamuy et al. 2000; Gallagher et al. 2005). This environmental effect and observationsof the local supernovae rate as a function of host galaxy properties (Cappellaro et al. 1999;Mannucci et al. 2005), motivated Scannapieco & Bildsten (2005) to parametrize the delaytime distribution, time between star formation and the appearance of SNe, and the rateswith a two-component model having a piece proportional to the SFR of the host galaxy (or prompt , they explode ∼ few × yr after an episode of star formation), and a second pieceproportional to the total stellar mass ( delayed component, they explode on scales of a fewGyr after the onset of star formation). The difference in age of the stellar populations of thesesubclasses suggests that the progenitors may also be different: prompt SN Ia would comefrom more massive progenitors. The host galaxies of all five SN 2002ic-like events knownare broadly consistent with the properties of the hosts of prompt
SN Ia, which suggest areal association given that the best studied SN IIa to date, SN 2002ic and SN 2005gj, havespectral characteristics similar to 1991T-like events.Several progenitors have been discussed in the literature for SN 2002ic and SN 2005gj.Livio & Riess (2003) proposed that SN 2002ic is a rare case of a double-degenerate binarysystem, a white dwarf (WD) and the core of an AGB star spiraling-in through gravitationalwave losses, in which the explosion occurs during or immediately after the common-envelopephase (a few hundred to a few thousand years of duration). The difference in line strengthsof the Balmer emission lines observed for SN 2002ic and SN 2005gj makes this scenariounlikely. Also, as Aldering et al. (2006) points out, in both SN 2002ic and SN 2005gj themass loss stopped only a few years before explosion, which is too short compared with the 30 –timescale for gravitational wave radiation to produce the merger of the core and the WD.Another possible progenitor initially proposed by Hamuy et al. (2003) and favored bythe models of Chugai et al. (2004), is the explosion of the Chandrasekhar-mass Carbon-Oxygen core of a massive AGB star in a degenerate medium, a supernova Type 1.5 (Iben & Renzini1983), where the dense Hydrogen-rich CSM would come from the outer layers of the AGB.In order for the core to grow to the Chandrasekhar mass, the radiatively driven winds fromthe AGB have to be weak enough, a condition that is only met in a very low-metallicityenvironment like the Galactic halo (Zijlstra 2004). At face value, the range of host galaxymetallicities for SN 2002ic-like events inferred from the luminosity-metallicity relation doesnot support the SN 1.5 scenario, although admittedly these are average metallicities and donot tell us the actual range of metallicities of the progenitors.Han & Podsiadlowski (2006) proposed that SN 2002ic could be produced through the“super-soft channel”, the most common single-degenerate model for the progenitors of SN Ia.In this scenario the white dwarf is accreting material from a main sequence, or slightlyevolved, relatively massive companion ( ∼ ⊙ ) and experiences a delayed dynamical in-stability that leads to a large amount of mass-loss from the system in the last few × yrbefore the explosion. Aldering et al. (2006) notes that the estimated main-sequence mass ofthe progenitor of SN 2005gj of ∼ ⊙ , calculated using the age of the starburst of its hostgalaxy, is consistent with the Han & Podsiadlowski (2006) model. Also, the predicted frac-tion of SN Ia that would be produced through the “delayed dynamical” channel is 0.1-1%,consistent with the limits we have obtain from the detection of SN 2005gj in the SDSS-IISurvey.In general terms, the progenitor model proposed by Han & Podsiadlowski (2006) suc-cessfully reproduces the observational properties of SN 2002ic and SN 2005gj. However, itis still very early in the study of this new sub-class of SN Ia. It would be interesting tosee in the near future the results of theoretical modeling exploring other single degenerateconfigurations (e.g., AGB donor) and detailed hydrodynamical modeling of the ejecta-CSMinteraction of SN 2005gj using the observations of the early photometric and spectroscopicevolution presented in this work.We are grateful for the assistance of the staffs at the many observatories (APO, MDM,LCO, ESO, La Palma) where data for this paper were obtained. We would like to thankK. Z. Stanek and R. Pogee for helpful discussions, and L. Watson for carefully reading anearlier version of this paper. We also thank S. Taubenberger for making the electronic formof the SN 2004aw data available, and S. Blondin for allowing us to use the SNID code before 31 –its public release. We wish to thank H. Tananbaum and the Chandra Observatory for thegenerous allotment of Director’s Discretionary Time. VVD research is supported by award REFERENCES
Aldering, G., et al. 2006, ArXiv Astrophysics e-printsAllen, D. A., et al. 1991, MNRAS, 248, 528Aretxaga, I., et al. 1999, MNRAS, 309, 343Arnett, W. D. 1982, ApJ, 253, 785Astier, P., et al. 2006, A&A, 447, 31Barentine, J., et al. 2005, Central Bureau Electronic Telegrams, 247, 1 (2005). Edited byGreen, D. W. E., 247Benetti, S., et al. 2006, ApJ, 653, L129Blondin, S. & Tonry, J. L. 2007, ApJ, in pressBranch, D., et al. 1981, ApJ, 244, 780Branch, D., et al. 2003, AJ, 126, 1489Branch, D., et al. 2000, PASP, 112, 217Branch, D., Romanishin, W., & Baron, E. 1996, ApJ, 465, 73Cappellaro, E., Evans, R., & Turatto, M. 1999, A&A, 351, 459Cardelli, J. A., Clayton, G. C., & Mathis, J. S. 1989, ApJ, 345, 245Chevalier, R. A. 1982, ApJ, 258, 790Chevalier, R. A. & Fransson, C. 1994, ApJ, 420, 268Chevalier, R. A. & Fransson, C. 2003, in LNP Vol. 598: Supernovae and Gamma-RayBursters, ed. K. Weiler, 171–194Chugai, N. N. 1997a, Ap&SS, 252, 225—. 1997b, Astronomy Reports, 41, 672Chugai, N. N., Chevalier, R. A., & Lundqvist, P. 2004, MNRAS, 355, 627Chugai, N. N. & Danziger, I. J. 1994, MNRAS, 268, 173Contardo, G., Leibundgut, B., & Vacca, W. D. 2000, A&A, 359, 876 33 –Delahaye, F. & Pinsonneault, M. H. 2006, ApJ, 649, 529Deng, J., Kawabata, K. S., Ohyama, Y., et al. 2004, ApJ, 605, L37Di Carlo, E., et al. 2002, ApJ, 573, 144Dilday, B., et. al. 2007, in preparationDraine, B. T. & McKee, C. F. 1993, ARA&A, 31, 373Dwarkadas, V. V. & Chevalier, R. A. 1998, ApJ, 497, 807Fisher, A., et al. 1999, MNRAS, 304, 67Fisher, A., et al. 1997, ApJ, 481, L89+Fisher, A. K. 2000, PhD thesis, AA(THE UNIVERSITY OF OKLAHOMA)Fransson, C., Lundqvist, P., & Chevalier, R. A. 1996, ApJ, 461, 993Freedman, W. L., et al. 2001, ApJ, 553, 47Frieman, J. A., et. al. 2007, submitted to AJFukugita, M., et al. 1996, AJ, 111, 1748Gal-Yam, A., et al. 2006, ArXiv Astrophysics e-printsGal-Yam, A., Ofek, E. O., & Shemmer, O. 2002, MNRAS, 332, L73Gallagher, J. S., et al. 2005, ApJ, 634, 210Garavini, G., et al. 2004, AJ, 128, 387Garnavich, P. M., et al. 2004, ApJ, 613, 1120Germany, L. M., et al. 2000, ApJ, 533, 320Gunn, J. E., et al. 1998, AJ, 116, 3040Gunn, J. E., et al. 2006, AJ, 131, 2332Hamuy, M., et al. 1993, PASP, 105, 787Hamuy, M., et al. 1995, AJ, 109, 1Hamuy, M., et al. 1996a, AJ, 112, 2391 34 –—. 1996b, AJ, 112, 2398Hamuy et al. 2000, AJ, 120, 1479Hamuy, M., et al. 2002, AJ, 124, 417Hamuy, M., et al. 2003, Nature, 424Hamuy, M., et al. 2006, PASP, 118, 2Han, Z. & Podsiadlowski, P. 2006, MNRAS, 368, 1095Hogg, D. W., et al. 2001, AJ, 122, 2129Holtzman, J., et. al. 2007, in preparationIben, Jr., I. & Renzini, A. 1983, ARA&A, 21, 271Immler et al. 2006, ApJ, 648, L119Immler, S., Petre, R., & Brown, P. 2005, IAU Circ., 8633, 2Ivezi´c, ˇZ., et al. 2007, ArXiv Astrophysics e-printsIvezi´c, ˇZ., et al. 2004, Astronomische Nachrichten, 325, 583Iwamoto, K., et al. 2000, ApJ, 534, 660Jeffery, D. J., et al. 2006, ArXiv Astrophysics e-printsJeffery, D. J., et al. 1992, ApJ, 397, 304Jha, S., et al. 1999, ApJS, 125, 73Kirshner, R. P., et al. 1993, ApJ, 415, 589Klein, R. I., McKee, C. F., & Colella, P. 1994, ApJ, 420, 213Kochanek, C. S., et al. 2001, ApJ, 560, 566Kotak, R., et al. 2004, MNRAS, 354, L13Kraft, R. P., Burrows, D. N., & Nousek, J. A. 1991, ApJ, 374, 344Krisciunas, K., Phillips, M. M., & Suntzeff, N. B. 2004, ApJ, 602, L81Leibundgut, B., et al. 1991, ApJ, 371, L23 35 –Leibundgut, B., et al. 1993, AJ, 105, 301Livio, M. & Riess, A. G. 2003, ApJ, 594, L93Lupton, R. H., Gunn, J. E., & Szalay, A. S. 1999, AJ, 118, 1406Mannucci, F., et al. 2005, A&A, 433, 807Martini, P., et al. 2004, in Ground-based Instrumentation for Astronomy. Edited by Alan F.M. Moorwood and Iye Masanori. Proceedings of the SPIE, Volume 5492, pp. 1653-1660 (2004)., ed. A. F. M. Moorwood & M. Iye, 1653–1660Matheson, T., et al. 2005, AJ, 129, 2352Mazzali, P. A., Danziger, I. J., & Turatto, M. 1995, A&A, 297, 509Miknaitis, G., et al.(2007), ArXiv Astrophysics e-printsMillard, J., et al. 1999, ApJ, 527, 746Modjaz, M., et al. 2007, ArXiv Astrophysics e-printsModjaz, M., et al. 2006, ApJ, 645, L21Morgan, C. W., et al. 2005, AJ, 129, 2504Nomoto, K., Thielemann, F.-K., & Yokoi, K. 1984, ApJ, 286, 644Nomoto, K., et al. 2005, in ASP Conf. Ser. 342: 1604-2004: Supernovae as CosmologicalLighthouses, ed. M. Turatto, S. Benetti, L. Zampieri, & W. Shea, 105–+Nugent, P., Kim, A., & Perlmutter, S. 2002, PASP, 114, 803Oke, J. B., & Gunn, J. E. 1983, ApJ, 266, 713Osterbrock, D. E. 1989, Astrophysics of gaseous nebulae and active galactic nuclei (Researchsupported by the University of California, John Simon Guggenheim Memorial Foun-dation, University of Minnesota, et al. Mill Valley, CA, University Science Books,1989, 422 p.Pastorello, A., et al. 2002, MNRAS, 333, 27Patat, F., et al. 2001, ApJ, 555, 900Perlmutter, S., et al. 1999, ApJ, 517, 565 36 –Persson, S. E., et al. 2002, AJ, 124, 619Phillips, M. M., et al. 1992, AJ, 103, 1632Phillips, M. M. 1993, ApJ, 413, L105Phillips , M. M., et al. 2007, PASP, 119, 360Pier, J. R., et al. 2003, AJ, 125, 1559Prieto, J., et al. 2005, IAU Circ., 8633Riess, A. G., et al. 1998, AJ, 116, 1009Riess, A. G., et al. 2004, ApJ, 607, 665Riess, A. G., et al. 2005, ApJ, 627, 579Riess, A. G., et al. 2006, ArXiv Astrophysics e-printsRigon, L., et al. 2003, MNRAS, 340, 191Salzer, J. J., et al. 2005, ApJ, 624, 661Scannapieco, E. & Bildsten, L. 2005, ApJ, 629, L85Schlegel, E. M. 1990, MNRAS, 244Schlegel, D. J., Finkbeiner, D. P., & Davis, M. 1998, ApJ, 500, 525Schmidt, B. P., et al. 1994, ApJ, 434, L19Smith, J. A., et al. 2002, AJ, 123Soderberg, A. M. & Frail, D. A. 2005, The Astronomer’s Telegram, 663, 1Spergel, D. N., et al. 2006, ArXiv Astrophysics e-printsStanek, K. Z., et al. 2006, Acta Astronomica, 56, 333Stritzinger, M., et al. 2006, A&A, 460, 793Taubenberger, S., et al.(2006), MNRAS, 371, 1459Tucker, D. L., et al. 2006, Astronomische Nachrichten, 327, 821Turatto, M., et al. 1993, MNRAS, 262, 128 37 –Turatto, M., et al. 2000, ApJ, 534, L57van Zee, L., Skillman, E. D., & Haynes, M. P. 2006, ApJ, 637, 269Wang, L., et al. 2004, ApJ, 604, L53Wood-Vasey, W. M., et al. 2007, ArXiv Astrophysics e-printsWood-Vasey, W. M., Wang, L., & Aldering, G. 2004, ApJ, 616, 339Wyder, T. K., et al. 2005, ApJ, 619, L15York, D. G., et al. 2000, AJ, 120, 1579Zijlstra, A. A. 2004, MNRAS, 348, L23
This preprint was prepared with the AAS L A TEX macros v5.2.
Table 1. SDSS ugriz and CSP u ′ g ′ r ′ i ′ photometry of comparison stars in common in the field of SN 2005gj. Star u g r i z ID α (J2000.0) δ (J2000.0) SDSS CSP SDSS CSP SDSS CSP SDSS CSP SDSS1 03 01 09.56 –00 33 52.5 18.639(027) 18.690(054) 17.367(018) 17.362(009) 16.837(022) 16.823(013) 16.643(018) 16.586(010) 16.549(020)2 03 01 06.29 –00 32 59.0 · · · · · · · · · · · · · · · · · · · · · · · · · · ·
39 –Table 2. SDSS and MDM ugriz photometry of SN 2005gj
JD Epoch a − , ,
000 (days) u g r i z
Source616.94 . . . ± ± ± ± ± . . . ± ± ± ± ± . . . ± ± ± ± ± . . . ± ± ± ± ± a Rest-frame days since the time of explosion (JD 2,453,637.93).
40 –Table 3. CSP u ′ g ′ r ′ i ′ photometry of SN 2005gj JD Epoch − , ,
000 (days) u ′ g ′ r ′ i ′ · · · · · · · · · · · · · · · · · · · · · · · · · · ·
41 –Table 4. CSP
Y J HK s photometry of SN 2005gj JD Epoch − , ,
000 (days)
Y J H K s Instrument700.71 59.1 16.565(015) 16.484(034) 16.253(030) · · ·
Retrocam704.68 62.9 16.591(015) 16.537(020) 16.315(033) · · ·
Retrocam709.66 67.6 16.628(015) 16.550(023) 16.271(029) · · ·
Retrocam714.58 72.2 16.673(015) 16.594(020) 16.389(028) · · ·
Retrocam718.65 76.0 16.725(016) 16.658(022) 16.364(037) · · ·
Retrocam722.62 79.8 16.832(016) 16.725(032) 16.490(017) 16.384(096) WIRC724.68 81.7 16.781(016) 16.716(028) · · · · · ·
Retrocam727.69 84.6 16.872(024) 16.757(036) · · · · · ·
Retrocam732.71 89.3 16.920(019) 16.839(025) · · · · · ·
Retrocam750.61 106.1 17.307(016) 17.152(016) 16.812(016) · · ·
WIRC755.57 110.8 17.337(024) 17.263(024) 16.891(025) 16.745(039) PANIC756.59 111.8 17.428(016) 17.269(016) 16.928(017) · · ·
WIRC757.60 112.7 17.423(024) 17.270(024) 16.938(025) 16.807(042) PANIC773.57 127.8 17.677(016) 17.476(016) · · · · · ·
WIRC776.54 130.6 17.538(026) 17.399(070) 17.081(068) · · ·
Retrocam777.55 131.5 17.648(038) 17.417(056) 17.129(119) · · ·
Retrocam782.55 136.2 17.513(039) 17.436(073) 17.072(114) · · ·
Retrocam783.55 137.2 17.766(024) 17.586(024) 17.271(025) 17.045(042) PANIC785.53 139.0 17.765(025) 17.570(022) 17.192(049) 17.152(177) WIRC788.54 141.9 17.869(016) 17.594(016) 17.335(019) · · ·
WIRC797.52 150.3 17.855(061) 17.563(118) 17.313(158) · · ·
Retrocam800.50 153.1 17.840(041) · · · · · · · · ·
Retrocam808.52 160.7 17.871(058) 17.551(079) · · · · · ·
Retrocam814.49 166.3 17.864(084) · · · · · · · · ·
RetrocamNote. — Uncertainties given in parentheses in thousandths of a magnitude.
42 –Table 5. Light-curve parameters for SN 2005gj
Parameter ValueTime of explosion a ± u max ± g max ± r max ± i max ± u max ± g max ± r max ± i max ± u max b ± g max ± r max ± i max ± M g,max –20.21 E ( B − V ) Gal c ± A g (Gal) 0.45 ± a JD-2 , , b Magnitudes at maximum in therest-frame, they have been cor-rected by Galactic extinction and K -corrections. We assume a negligi-ble extinction in the host galaxy. c From Schlegel et al. (1998)
43 –Table 6. Spectroscopic observations of SN 2005gj
JD Epoch Wavelength Resolution a Exposure − , ,
000 (days) Instrument Range (˚A) (˚A) (s)644.92 6.6 MDM-CCDS 3850 – 7308 15 1200646.95 8.5 ARC-DIS 3824 – 10192 7 1800650.84 12.2 ARC-DIS 3600 – 9597 7 1000655.87 16.9 MDM-CCDS 3823 – 7283 15 1800665.92 26.4 MDM-CCDS 3883 – 7341 15 2700668.83 29.1 MDM-CCDS 3886 – 7346 15 2700676.79 36.6 MDM-CCDS 3882 – 7338 15 3600684.73 44.1 WHT-ISIS 3924 – 8901 4 1800686.79 46.0 MDM-CCDS 3858 – 7315 15 2700698.67 57.2 duPont-ModSpec 3780 – 7290 7 2700699.67 58.2 duPont-ModSpec 3780 – 7290 7 2700700.76 59.2 MDM-CCDS 3933 – 7391 15 2700702.73 61.1 MDM-CCDS 3856 – 7310 15 2700712.73 70.5 MDM-CCDS 3831 – 7286 15 2700722.71 79.9 NTT-EMMI 4000 – 10200 9 2700724.66 81.7 duPont-WFCCD 3800 – 9235 6 2700725.65 82.6 duPont-WFCCD 3800 – 9235 6 2700726.66 83.6 duPont-WFCCD 3800 – 9235 6 3600727.67 84.5 duPont-WFCCD 3800 – 9235 6 3600728.67 85.5 duPont-WFCCD 3800 – 9235 6 3600729.67 86.4 MDM-CCDS 3915 – 7373 15 2700737.70 94.0 MDM-CCDS 3909 – 7368 15 2700751.60 107.1 NTT-EMMI 3200 – 10200 9 2700755.62 110.9 MDM-CCDS 3844 – 7299 15 3600759.61 114.6 Magellan-LDSS-3 3788 – 9980 3 3600799.52 152.2 duPont-WFCCD 3800 – 9235 6 1200Note. — Most of the spectra are the combination of multiple observation, the totalexposure is given. a Average resolution obtained from the FWHM of arc-lamp lines.
44 –Table 7. K -corrections of SN 2005gj Epoch(days) K u K g K r K i K z
45 –Table 8. Derived integrated luminosity and black-body fits.
Epoch log L ( u → i ) log L ( u → z ) log L bol T bb a R bb b (days) (erg s − ) (erg s − ) (erg s − ) (K) (10 cm) χ ν c a Black-body temperature from the fits to the broadband photometry; 1 σ uncertainty aregiven in parenthesis. b Black-body radius from the fits to the broadband photometry; 1 σ uncertainty are givenin parenthesis and include a 10% uncertainty in the distance to SN 2005gj. c χ per degree of freedom of the black-body fits.
46 –Table 9. Library of spectra used in SNID
SN Name Class Epochs Reference1990N Ia normal -14, -13, -8, -7, -6, 0, 4, 8, 15, 18, 39 11991T Ia 91T -9, -8, -7, -6, -5, -4, -2, -1, 9, 10, 11, 12, 15, 16, 17, 18, 19, 20, 21, 22, 23, 27, 43, 44, 47, 48, 51, 69, 77 2, 31991bg Ia 91bg 1, 3, 16, 18, 25, 32, 33, 46, 54, 85 41992A Ia normal -5, -1, 3, 5, 6, 7, 9, 11, 16, 17, 24, 28 51994I Ic normal -6, -4, -3, 0, 1, 2, 21, 22, 23, 24, 26, 30, 36, 38 61997ef Ic broad -14, -12, -11, -10, -9, -6, -5, -4, 7, 13, 14, 16, 17, 19, 22, 24, 27, 41, 45, 47, 49, 75, 80, 81 71998aq Ia normal -9, -8, -3, 0, 1, 2, 3, 4, 5, 6, 7, 19, 21, 24, 31, 32, 36, 51, 55, 58, 60, 63, 66, 79, 82, 91, 211, 231, 241 81998bu Ia normal -3, -2, -1, 9, 10, 11, 12, 13, 14, 28, 29, 30, 31, 32, 33, 34, 35, 36, 37, 38, 39, 40, 41, 42, 43, 44, 57 91998bw Ic broad -9, -7, -6, -3, -2, -1, 1, 3, 4, 6, 9, 11, 12, 13, 19, 22, 29, 45, 64, 125, 200, 337, 376 101999aa Ia 91T -11, -7, -3, -1, 5, 6, 14, 19, 25, 28, 33, 49, 47, 51 111999ee Ia normal -9, -7, -2, 0, 3, 8, 10, 12, 17, 20, 23, 28, 33, 42 121999ex Ic normal -1, 4, 13 121999by Ia 91bg -4, -3, -2, -1, 2, 3, 5, 6, 7, 8, 10, 11, 25, 29, 31, 33, 42 132002ap Ic broad -5, -4, 3, 8, 10, 17, 19 142004aw Ic normal 1, 5, 6, 8, 15, 21, 22, 26, 28, 39, 35, 44, 49, 63, 64, 236, 260, 413 152006aj Ic broad -6, -5, -4, -3, -2, -1, 0, 2, 3 16References. — (1) Leibundgut et al. (1991); (2) Jeffery et al. (1992); (3) Schmidt et al. (1994); (4) Leibundgut et al. (1993); (5) Kirshner et al.(1993); (6) Millard et al. (1999); (7) Iwamoto et al. (2000); (8) Branch et al. (2003); (9) Jha et al. (1999); (10) Patat et al. (2001); (11)Garavini et al. (2004); (12) Hamuy et al. (2002); (13) Garnavich et al. (2004); (14) Gal-Yam et al. (2002); (15) Taubenberger et al. (2006);(16) Modjaz et al. (2006).
Table 10. Results of the Gaussian fits to H α and H β features H α (narrow) H α (broad) H β JD Epoch FWHM a flux b FWHM a flux b FWHM a flux b − , ,
000 (days) (km s − ) (10 − erg s − cm − ) (km s − ) (10 − erg s − cm − ) (km s − ) (10 − erg s − cm − )644.92 6.6 . . . 0.24(0.02) 1575 0.58(0.06) 776 0.50(0.05)646.95 8.5 137 0.37(0.04) 1481 0.83(0.08) 1307 0.52(0.06)650.84 12.2 314 0.69(0.07) 1731 1.25(0.13) 1462 0.75(0.10)655.87 16.9 . . . 0.57(0.06) 1555 1.11(0.11) 1339 0.75(0.09)665.92 26.4 . . . 0.53(0.05) 1569 1.03(0.10) 523 0.48(0.06)668.83 29.1 . . . 0.37(0.04) 1234 0.99(0.10) 1275 0.55(0.07)676.79 36.6 . . . 0.45(0.05) 1513 0.77(0.08) 884 0.28(0.04)686.79 46.0 . . . 0.41(0.04) 1836 0.71(0.07) . . . 0.15(0.02)698.67 57.2 . . . 0.45(0.04) 2115 0.91(0.10) . . . 0.15(0.03)699.67 58.2 . . . 0.34(0.04) 2053 0.76(0.08) . . . 0.14(0.03)700.76 59.2 . . . 0.36(0.04) 1830 0.77(0.08) 620 0.16(0.02)702.73 61.0 . . . 0.41(0.04) 1978 0.67(0.07) . . . 0.11(0.02)712.73 70.5 . . . 0.41(0.04) 2357 1.02(0.11) . . . 0.11(0.02)722.71 79.9 . . . 0.34(0.04) 2413 1.03(0.11) 490 0.12(0.03)724.66 81.7 . . . 0.34(0.04) 2137 0.99(0.10) 1067 0.19(0.03)725.65 82.6 . . . 0.34(0.03) 2260 1.02(0.11) 568 0.11(0.02)726.66 83.6 . . . 0.35(0.04) 2322 1.10(0.11) . . . 0.08(0.02)727.67 84.5 160 0.32(0.03) 2364 1.02(0.10) . . . 0.10(0.02)728.67 85.5 . . . 0.20(0.02) 1802 1.24(0.14) 1127 0.19(0.03)729.67 86.4 . . . 0.37(0.04) 2687 1.15(0.12) 680 0.14(0.02)737.70 94.0 . . . 0.29(0.03) 1941 0.85(0.09) 459 0.09(0.01)755.62 110.9 . . . 0.32(0.03) 2236 1.14(0.12) 1031 0.16(0.02)799.52 152.2 525 0.51(0.05) 3809 2.18(0.22) 1669 0.23(0.03) a FWHM is not presented when the spectral resolution is bigger than the measured value. b σ uncertainties are given in parentheses.
48 –
SN1615 141312 1110 98 76 543 21 NE Fig. 1.— r ′ -band image (3 . ′ × . ′ ) of the field around SN 2005gj obtained with the Swope-1m telescope at LCO. North is up and east is to the left. Sixteen comparison stars in commonbetween SDSS and CSP used to derive differential photometry of the SN are labeled as inTable 1. The SN is close to the center of the field. 49 –Fig. 2.— Spectra of SN 2005gj obtained from Oct. 1 ( ∼ ∼
61 days after explosion) of 2005. The sequence show the dramatic spectral evolution ofthe SN from a very blue continuum with strong Hydrogen-Balmer lines in emission in theearly phases, resembling the spectrum of a Type IIn SN, to a Type Ia supernova-dominatedcontinuum with broad absorption and emission features (P-cygni profiles) of blended Fe II and Fe III profiles. The spectra are shown in logarithmic flux scale and a constant shifthas been applied for clarity. The wavelength is in the rest-frame corrected using z = 0 . ∼
71 days afterexplosion) and Mar. 6, 2006 ( ∼
152 days after explosion). The labels, axis and symbols arethe same as in Figure 2. The earth symbol shows the position of a telluric feature present insome of the spectra. 51 –Fig. 4.— Observed light curves of SN 2005gj from SDSS ( open circles ), MDM ( open squares )and CSP/Swope ( filled triangles ). The error bars are smaller than the symbols. For clarity,the light curves have been shifted by an arbitrary constant. 52 –Fig. 5.— Difference between the synthetic g − r color calculated from the spectra and theobserved color from the photometry. We do not include the latest spectra obtained on Jan24. and Mar. 6 because there is no contemporaneous photometric data. 53 –Fig. 6.— Time evolution of the colors of SN 2005gj ( filled circles ). For comparison we alsoshow the color evolution of the overluminous Type Ia SN 1991T ( solid line ), the Type IInSN 1999el ( dashed line ), and two previous cases of Type Ia strongly interacting with itscircumstellar medium, SN 2002ic ( stars ) and SN 1997cy ( triangles ). 54 –Fig. 7.— Absolute ugri light curves of SN 2005gj ( filled circles ). For comparison we alsoshow the absolute light curves of the overluminous Type Ia SN 1991T ( solid line ), SN 2002ic( stars ) and SN 1997cy ( triangles ). The error bar in the lower right pannel represents thetypical error in the absolute magnitudes dominated by a 10% uncertainty in the Hubbleconstant. 55 –Fig. 8.— Absolute light curves of SN 2005gj in the NIR: J ( top panel ), H ( middle panel )and K s ( bottom panel ). For comparison we also show the absolute light curves of a normalType Ia ( solid line ) and the Type IIn SN 1999el ( open triangles ). The error bar in the lowerright of each panel represents the typical error in the absolute magnitudes dominated by a10% uncertainty in the Hubble constant. 56 –Fig. 9.— Examples of black-body fits ( solid line ) to the SED of SN 2005gj obtained by trans-forming the rest-frame ugriz magnitudes to monochromatic fluxes at the effective wavelengthof the filters ( filled circles ). These examples show the quality (i.e., goodness-of-fit) range ofthe black body fits at different epochs: χ ν =0.3 ( left panel ), 2.4 ( middle ), 4.4 ( right ). Theunits of flux density in the y-axis are mJy = 10 − erg s − cm − Hz − . 57 –Fig. 10.— Top panel:
Quasi-bolometric ( open circles ) and bolometric light curves ofSN 2005gj ( filled circles ). The bolometric light luminosities were obtained after applyingbolometric corrections calculated from black-body fits to the optical SED obtained from the ugriz photometry. The dashed line shows the best-fit linear decay of 0 .
014 mag day − . The middle and bottom panels show the evolution of the black-body temperature and radius,respectively. 58 –Fig. 11.— Comparison of spectra of SN 2005gj at 26, 46, 62 and 83 days after explosionwith comparable epoch spectra of SN 2002ic (Hamuy et al. 2003) and SN 1997cy (fromSUSPECT database). The spectra are plotted on a logarithmic flux scale and shifted by anarbitrary constant. The wavelength was shifted to the restframe using z = 0 . blue line ); (2) fourth order polynomial ( green line ). The results ofthe fits are in red and the spectra of SN 2005gj, corrected by Galactic extinction in the lineof sight are in black . The epochs of the spectra are shown in the upper right of each panel. 61 –Fig. 14.— Results from the Gaussian fits to H α and H β emission features as a functionof time. Top left panel:
FWHM of the H α -broad and H β . Top right panel:
Luminosity ofthe narrow and broad Gaussian components of H α . Bottom left panel:
Luminosity of H β . Bottom right panel:
Balmer decrement, ratio of total fluxes in H α and H β lines. 62 –Fig. 15.— Line profiles of H β ( top ) and H α ( bottom ) in the two highest resolution spectraof SN 2005gj obtained with WHT+ISIS on day 44 ( left ) and Magellan+LDSS-3 on day115 ( right ). The features show clear P-cygni profiles with weak absorption minima at ∼−
200 kms − , demonstrating the presence of a slowly moving outflow. 63 –Fig. 16.— Identification of lines in the spectrum of SN 2005gj obtained at 17 days afterexplosion (2 days before g maximum). The red line shows the best fit synthetic spectrumgenerated with the SYNOW code. The lines of SN 2005gj are typical of SN 1991T aroundmaximum light ( blue line ), and very similar to the spectrum of SN 2002ic around maximum( green linegreen line