Supernova Collisionless Shocks

X-ray emission from supernovae in dense circumstellar matter environments: A search for collisionless shocks

Abstract

The optical light curve of some supernovae (SNe) may be powered by the outward diffusion of the energy deposited by the explosion shock (so-called shock breakout) in optically thick () circumstellar matter (CSM). Recently, it was shown that the radiation-mediated and -dominated shock in an optically thick wind must transform into a collisionless shock and can produce hard X-rays. The X-rays are expected to peak at late times, relative to maximum visible light. Here we report on a search, using Swift-XRT and Chandra, for X-ray emission from 28 SNe that belong to classes whose progenitors are suspected to be embedded in dense CSM. Our sample includes 19 type-IIn SNe, one type-Ibn SN and eight hydrogen-poor super-luminous SNe (SLSN-I; SN 2005ap like). Two SNe (SN 2006jc and SN 2010jl) have X-ray properties that are roughly consistent with the expectation for X-rays from a collisionless shock in optically thick CSM. Therefore, we suggest that their optical light curves are powered by shock breakout in CSM. We show that two other events (SN 2010al and SN 2011ht) were too X-ray bright during the SN maximum optical light to be explained by the shock breakout model. We conclude that the light curves of some, but not all, type-IIn/Ibn SNe are powered by shock breakout in CSM. For the rest of the SNe in our sample, including all the SLSN-I events, our X-ray limits are not deep enough and were typically obtained at too early times (i.e., near the SN maximum light) to conclude about their nature. Late time X-ray observations are required in order to further test if these SNe are indeed embedded in dense CSM. We review the conditions required for a shock breakout in a wind profile. We argue that the time scale, relative to maximum light, for the SN to peak in X-rays is a probe of the column density and the density profile above the shock region. The optical light curves of SNe, for which the X-ray emission peaks at late times, are likely powered by the diffusion of shock energy from a dense CSM. We note that if the CSM density profile falls faster than a constant-rate wind density profile, then X-rays may escape at earlier times than estimated for the wind profile case. Furthermore, if the CSM have a region in which the density profile is very steep, relative to a steady wind density profile, or the CSM is neutral, then the radio free-free absorption may be low enough, and radio emission may be detected.

Subject headings:
stars: mass-loss — supernovae: general — supernovae: individual
26

1. Introduction

Circumstellar Matter (CSM) around supernova (SN) progenitors may play an important role in the emission and propagation of energy from SN explosions. The interaction of the SN radiation with optically thin CSM shells may generate emission lines, with widths that are representative of the shell velocity (i.e., type-IIn SNe; Schlegel 1990; Kiewe et al. 2012). The interaction of SN ejecta with the CSM can power the light curves of SNe by transformation of the SN kinetic energy into photons. In cases where considerable amounts of optically thin (and ionized) material is present around the exploding star, synchrotron and free-free radiation can emerge, and inverse Compton scattering can generate X-ray photons (e.g., Chevalier & Fransson 1994; Horesh et al. 2012; Krauss et al. 2012).

For the type-IIn SN PTF 09uj, Ofek et al. (2010) suggested that a shock breakout can take place in optically-thick wind (see also Grassberg, Imshennik, & Nadyozhin 1971; Falk & Arnett 1977; Chevalier & Irwin 2011; Balberg & Loeb 2011). This will happen if the Thomson optical depth within the wind profile is , where is the speed of light, and is the shock speed. Ofek et al. (2010) showed that shock breakout in wind environments produces optical displays that are brighter and have longer time scales than those from surfaces of red supergiants (e.g., Colgate 1974; Matzner & McKee 1999; Nakar & Sari 2010; Rabinak & Waxman 2011; Couch et al. 2011). Chevalier & Irwin (2011) extended this picture. Specifically, they discussed CSM with a wind profile in which the wind has a cutoff at a distance . If the optical depth at is then the light curve of the supernova will have a slow decay (e.g., SN 2006gy; Ofek et al. 2007; Smith et al. 2007). If the optical depth at is , then it will have a faster decay (e.g., SN 2010gx; Pastorello et al. 2010a; Quimby et al. 2011a). Moriya & Tominaga (2012) investigated shock breakouts in general wind density profiles of the form . They suggested that, depending on the power-law index , shock breakouts in wind environments can produce bright SNe without narrow emission lines (e.g., SN 2008es; Gezari et al. 2009; Miller et al. 2009).

Recently Katz et al. (2011) and Murase et al. (2011) showed that if the progenitor is surrounded by optically thick CSM then a collisionless shock is necessarily formed during the shock breakout. Moreover, they argued that the energy emitted from the collisionless shock in the form of high-energy photons and particles is comparable to the shock breakout energy. Furthermore, this process may generate high-energy ( TeV) neutrinos. Although Katz et al. (2011) predicted that the photons are generated with energy typically above 60 keV, it is reasonable to assume that some photons will be emitted with lower energy. Chevalier & Irwin (2012) showed that Comptonization and inverse Compton of the high-energy photons is likely to play an important role, and that the high-energy photons will be absorbed.

Svirski, Nakar & Sari (2012) discuss the X-ray emission from collisionless shocks. They show that at early times the X-rays will be processed into the optical regime by the Compton process. Therefore, at early times, the optical emission will be about times stronger than the high-energy emission. With time, the X-ray emission will become stronger, while the optical emission will decay. They conclude that for a CSM with a steady wind profile (), X-ray emission may peak only at late times, roughly 10–50 times the shock breakout time scale. The shock breakout time scale, , is roughly given by the diffusion time scale at the time of shock breakout. This time scale is also equivalent to the radius at which the shock breaks out () divided by the shock velocity (; Weaver 1976). If the main source of optical photons is due to diffusion of the shock breakout energy, the SN optical light rise time, , will be equivalent to the shock breakout time scale. Therefore, X-ray flux measurements and spectra of SNe embedded in dense CSM starting from the explosion until months or years after maximum light are able to measure the properties of the CSM around the SN progenitors and the progenitor mass-loss history. This unique probe into the final stages of massive star evolution has been only partially exploited, at best.

Herein, we analyze the X-ray data for 28 SNe with light curves that may be powered by a shock breakout from dense CSM, and for which Swift-XRT (Gehrels et al. 2004) observations exist. We use this sample to search for X-ray signatures of collisionless shocks – emission at late times (months to years after peak optical luminosity). We suggest that these signals were observed in several cases, most notably in SN 2006jc (Immler et al. 2008) and SN 2010jl (Chandra et al. 2012a). Finally, we review the conditions for a shock breakout in CSM with a wind profile and discuss, the importance of bound-free absorption and the possibility to detect radio emission from such SNe.

The structure of this paper is as follows. in §2 we present the SN sample, while §3 presents the X-ray observations. We review and discuss the model in §4, and discuss the observations in context of the model in §5. Finally, we conclude in §6.

2. Sample

Our sample of SNe is based on SNe found by amateur astronomers, and several surveys, including the Lick Observatory Supernova Search (LOSS; Li et al. 2000), the Catalina Real-Time Transient Survey (CRTS; Drake et al. 2009a), Pan-STARRS1 (PS1; Kaiser et al. 2002), and the Palomar Transient Factory27 (PTF; Law et al. 2009; Rau et al. 2009). Two SNe, PTF 09drs and PTF 10tel, are reported here for the first time. We note that many of the nearby or luminous SNe found by PTF are also observed by Swift.

We selected a sample of SNe in which the main source of energy may be explained by diffusion of the explosion shock energy through optically thick CSM around the progenitor. First, we include type-IIn SNe within 200 Mpc. SNe that belong to this class show narrow hydrogen emission lines. This is an indication of the presence of optically thin material somewhere around the progenitor. However, it is unlikely that all SNe showing narrow hydrogen emission lines in their spectra are powered mainly by the diffusion of shock energy in an optically thick environment. One reason is that some type-IIn SNe show X-ray emission near maximum optical light, which is not expected when optically thick CSM is present (see §4). Furthermore, Moriya & Tominaga (2012) suggested that not all SNe powered by interaction of the ejecta with slow-moving material will necessarily have narrow emission lines in their spectrum. We note that some of the type-IIn SNe in our sample are peculiar (e.g., SN 2010jp/PTF 10aaxi; Smith et al. 2012).

Another relevant, but rare, class of objects are type-Ibn SNe. This class is defined by the lack of hydrogen lines and presence of narrow Helium emission lines. The only SN of this type in our sample is SN 2006jc (Nakano et al. 2006; Foley et al. 2007; Pastorello et al. 2008).

The third class of SNe we investigate here is the small group of hydrogen-poor super-luminous SN (SLSN-I; see review in Gal-Yam et al. 2012). Quimby et al. (2011a) used the spectra of several such events found by PTF, at intermediate redshift (), to show that these events, as well as SCP 06F6 (Barbary et al. 2009), and SN 2005ap (Quimby et al. 2007) are spectroscopically similar. This group of SNe continues to grow with new discoveries (e.g., Chomiuk et al. 2011; Leloudas et al. 2012), and their hosts were studied in Neill et al. (2011). Although the nature of these events is not understood (e.g., Kasen & Bildsten 2010), Quimby et al. (2011a) suggested that they may be powered by a pulsational pair-instability SN (Rakavy, Shaviv, & Zinamon 1967; Woosley, Blinnikov, & Heger 2007). According to this tentative model, the SN ejecta interacts with a dense shell of material, enriched with intermediate mass elements, that was ejected by the progenitor during previous explosions (see also Ginzburg & Balberg 2012). This model is tentatively supported by observations of SN 2006oz (Leloudas et al. 2012) that may show a dip in the light curve followed by re-brightening. Moriya & Maeda (2012) interpret the dip as an increase in the opacity due to ionization of the massive shell/CSM as it interact with the ejecta.

Our sample, presented in Table 1, consists of eight SLSN-I objects, 19 type-IIn SNe and a single type-Ibn SN. We note that the spectra of SNe with PTF names, as well as some other SNe, are available online from the WISeREP28 website (Yaron & Gal-Yam 2012).

Name Type z N
deg deg day mag MJD  cm erg s
PTF 09atu SLSN-I 247.602288 23.640289 30: 22.5 0.501 55060 4.05 0.7
PTF 09cnd SLSN-I 243.037252 51.487820 50 22.8 0.258 55080 1.67 0.02
PTF 09cwl/SN 2009jh SLSN-I 222.291998 29.419833 50 22.5: 0.349 55060 1.51 0.4
SN 2010gx/PTF 10cwr SLSN-I 171.444479 8.828099 20 21.7: 0.231 55280 3.78 0.07
PTF 10hgi SLSN-I 249.446009 6.208978 50 20.3 0.096 55370 6.06 0.1
PTF 11dij SLSN-I 207.740690 26.278562 40: 21.1: 0.143 55690 1.21 0.06
PTF 11rks SLSN-I 24.939618 29.924170 20 21.0 0.20 55945 5.27 0.07
PS1-12fo SLSN-I 146.553792 19.841306 21.0: 0.175 55956 2.79 0.2
SN 2006jc Ibn 139.366667 41.908889 17.8 0.006 54020 1.00 0.04
PTF 09drs IIn 226.625665 60.594271 40: 17.8: 0.045 55210 1.72 1.2
SN 2010jl/PTF 10aaxf IIn 145.722208 9.494944 20.6 0.011 55500 3.05 0.004
SN 2010jp/PTF 10aaxi IIn 94.127702 21.410012 14.6: 0.01 55520 11.0 0.06
SN 2010jj/PTF 10aazn IIn 31.717743 44.571558 18.0: 0.016 55530 9.38 0.03
SN 2010bq/PTF 10fjh IIn 251.730659 34.159641 18.5 0.032 55310 1.79 0.2
PTF 11iqb IIn 8.520150 9.704979 10: 18.4 0.013 55780 2.79 0.01
SN 2007bb IIn 105.281083 51.265917 17.6: 0.021 54192 7.04 0.09
SN 2007pk IIn 22.946125 33.615028 17.3: 0.017 54423 4.72 0.04
SN 2008cg IIn 238.563125 10.973611 19.4: 0.036 54583: 3.65 0.02
SN 2009au IIn 194.941667 29.602083 16.5: 0.009 54901: 6.42 0.03
SN 2010al IIn 123.566292 18.438389 16.0: 0.0075 55268: 3.92 0.3
SN 2011ht IIn 152.044125 51.849167 30 14.2: 0.004 55880 0.78 0.05
SN 2011hw IIn 336.560583 34.216417 19.1: 0.023 55883: 10.2 0.004
PTF 10tel IIn 260.377817 48.129834 17 18.5 0.035 55450 2.34 0.1
SN 2011iw IIn 353.700833 24.750444 18.1: 0.023 55895: 1.61 0.02
SN 2005db IIn 10.361625 25.497667 16.8: 0.0153 53570: 4.17 0.04
SN 2005av IIn 311.156583 68.752944 17.8: 0.0104 53453: 4.85 0.02
SN 2003lo IIn 54.271333 5.038139 15.8: 0.0079 53005: 4.87 0.01
SN 2002fj IIn 130.187917 4.127361 18.5 0.0145 52532 3.12 0.007

Note. – The sample of SNe with X-ray observations. Type refer to SN type, and are the J2000.0 right ascension and declination, respectively. is the approximate rise time of the SN optical light curve. The rise time is deduced from various sources including PTF and KAIT photometry and the literature listed in the references. We note that sign indicates the range rather than the uncertainty in . The column sign indicates uncertain value. is the approximate absolute -band magnitude at maximum light (ignoring -corrections). refers to the object redshift. If the galaxy is nearby and has a direct distance measurements in the NASA Extragalactic Database (NED), we replaced the observed redshift by the redshift corresponding to the luminosity distance of the galaxy. is the MJD of maximum light, and is the Galactic neutral hydrogen column density for the source position (Dickey & Lockman 1990). is the X-ray luminosity or the 2- upper limit on the X-ray luminosity in the 0.2–10 keV band. Finally, is the ratio between the X-ray measurments or limit () and the peak visible light luminosity.
References:
PTF 09atu: Quimby et al. (2011a).
PTF 09cnd: Quimby et al. (2011a); Chandra et al. (2009); Chandra et al. (2010).
PTF 09cwl: SN 2009jh; Quimby et al. (2011a); Drake et al. (2009b).
SN 2010gx: PTF 10cwr; Mahabal et al. (2010); Quimby et al. (2010a); Pastorello et al. (2010b); Quimby et al. (2011).
PTF 10hgi: Quimby et al. (2010b).
PTF 11dij: Drake et al. (2011a); Drake et al. (2011b); Quimby et al. (2011b).
PTF 11rks: Quimby et al. (2011c)
PS1-12fo: Drake et al. (2012); Smartt et al. (2012); Maragutti et al. (2012).
PTF 09drs: Reported here for the first time.
SN 2010jl: PTF 10aaxf; Newton et al. (2010); Stoll et al. (2011).
SN 2010jp: PTF 10aaxi; A peculiar type-IIn supernova; Maza et al. (2010); Challis et al. (2010b); Smith et al. (2012).
SN 2010jj: PTF 10aazn; Rich (2010b); Silverman et al. (2010b).
SN 2010bq: PTF 10fjh; Duszanowicz et al. (2010); Challis et al. (2010a).
PTF 11iqb: Parrent et al. (2011); Horesh et al. (2011).
SN 2007bb: Joubert & Li(2007); Blondin et al. (2007). and are based on unpublished KAIT photometry.
SN 2007pk: A peculiar type-IIn supernova (Parisky et al. 2007). and are based on unpublished KAIT photometry.
SN 2008cg: Drake et al. (2008); Blondin et al. (2008); Filippenko et al. (2008; cbet 1420); Spectrum is similar to SN 1997cy (Filippenko 2008).
SN 2009au: Pignata et al. (2009); Stritzinger et al. (2009).
SN 2010al: Spectrum is similar to SN 1983K with He ii, N iii, and H i emission lines; Rich et al. (2010a); Stritzinger et al. (2010); Silverman et al. (2010a).
SN 2011ht: Boles et al. (2011); Prieto et al. (2011); Roming et al (2012).
SN 2011hw: Dintinjana et al. (2011).
PTF 10tel: Reported here for the first time.
SN 2011iw: Mahabal et al. (2011).
SN 2005db: Blanc et al. (2005); Monard (2005b); Kiewe et al. 2012).
SN 2005av: Monard (2005a); Salvo et al. (2005).
SN 2003lo: Puckett et al. (2004); Matheson et al. (2004).
SN 2002fj: Monard & Africa (2002); Hamuy (2002).

Table 1SN sample

3. Observations

For each Swift-XRT image of an SN, we extracted the number of X-ray counts in the 0.2–10 keV band within an aperture of (3 pixels) radius centered on the SN position. We chose small aperture in order to minimize any host galaxy contamination. We note that this aperture contains % of the source flux (Moretti et al. 2004). The background count rates were estimated in annuli around each SN, with an inner (outer) radius of (). For each SN that has Swift-XRT observations, we searched for Chandra observations. The Chandra observations were analyzed in a similar manner with an extraction aperture of and background annuli with an inner (outer) radius of (). All the Swift-XRT X-ray measurements are listed in Table 2 (the full table is available in the electronic version). In addition, in Table 2, for each object we give the count rate in up to four super-epochs: (i) all the observations taken prior to maximum light, or discovery date if time of maximum light is not known; (ii) all the observations taken between maximum light and 300 days after maximum light; (iii) all the observations taken more than 300 days after maximum light; and (iv) all the observations at the position of the SN.

In each epoch, and super-epoch, we also provide the source false alarm probability (FAP), which is the probability that the X-ray counts are due to the X-ray background rather than a source. This probability is estimated as 1 minus the Poisson cumulative distribution to get a source count rate smaller than the observed count rate, assuming the expectancy value of the Poisson distribution equal to the background counts29 within an area equivalent to the aperture area. We note that in some cases X-ray emission from the host galaxy will tend to produce some seemingly significant, but actually ”false alarm” detections under these assumptions. In cases in which FAP, we estimated also the 2- upper limit on the count rate (Gehrels 1986). The count-rate measurements or upper limits are converted to luminosity in the 0.2–10 keV band using the PIMMS30 web tool and assuming that: the aperture in which we extracted the source photometry contains 37% of the source photons (Moretti et al. 2004); Galactic neutral hydrogen column density in the position of the sources as listed in Table 1 (Dickey & Lockman 1990); a spectrum of , where has units of photons cm s (motivated in §4); and the luminosity distance to the SNe calculated using the redshifts listed in Table 1 and  km s Mpc; ; ; (the 3rd year WMAPSNLS cosmology; Spergel et al. 2007).

Name Exp. CR CR CR CR FAP
start end
day day s ct ks ct ks ct ks ct ks erg s
PTF09atu 4858.0 0.62 1.00
PTF09cnd 3441.8 0.87 1.00
3557.2 0.84 1.00
3273.1 0.92 1.00
3997.8 0.75 1.00
2232.1 1.34 1.00
2980.6 1.01 1.00
2026.8 1.48 1.00
1910.6 1.57 1.00
16501.9 0.18 1.00
6917.9 0.43 1.00
23419.8 0.13 1.00

Note. – Summary of all the 305 Swift-XRT flux measurements of the 28 SNe in our sample. For each SN we list the observation date () relative to the listed in Table 1. Rows that have both a start day and end day listed, give ’super-epoch’ measurements as described in the main text. Exp. is the exposure time, CR, CR, CR are the source counts rate, lower error and upper error, respectively. CR is the 2- upper limit on the source count rate, FAP is the false alarm probability (see text), and is the source luminosity, or the the 2- upper limit on the luminosity, in the 0.2–10 keV band. If FAP than we provide the 2- upper limits, otherwise the measurements are given along with the errors. This table is published in its entirety in the electronic edition of the Astrophysical Journal. A portion of the full table is shown here for guidance regarding its form and content.

Table 2Swift-XRT X-ray measurements

Several objects listed in Table 2 show count rates which may deviate from zero. Here we discuss the observations of all seven sources that have FAP in at least one of the epochs or super-epochs. We note that Table 2 contains 305 epochs and super-epochs. Therefore, we expect about three random false detections. Interpretation of these observations is discussed in §5.

SN 2010jl/PTF 10aaxf (Fig. 1): This SN has a large number of Swift-XRT observations, as well as Chandra ACIS-S observations in five epochs (Chandra et al. 2012a), of which three are public (PIs: Pooley; Tremonti). The host, SDSS J094253.43092941.9, is an irregular star-forming galaxy. The Binned Swift-XRT and Chandra light curves, as well as the PTF -band light curve, are presented in Figure 1.

Figure 1.— The Swift (circles) and Chandra (squares) X-ray light curve extracted at the position of SN 2010jl. The unabsorbed flux was calculated using PIMMS assuming Galactic column density of  cm, and power-law spectrum. We note that the Chandra observations show a possible extended source near the SN location. This additional source may contaminate the Swift-XRT measurements and can explain the small discrepancy between Chandra and Swift-XRT. Alternatively, the discrepancy between the Chandra and Swift light curves can be explained if the X-ray spectrum is harder or the is larger than we assumed. We note that for N which is 1000 larger than the Galactic value, the unabsorbed Swift (Chandra) flux will be about 5.2 (7.2) times larger. For reference, the grey circles show the PTF -band luminosity of this SN scaled by 0.01. The PTF -band luminosity was calibrated using the method described in Ofek et al. (2012a) and calibration stars listed in Ofek et al. (2012b).

SN 2006jc (Fig. 2): this is the only SN in our sample that belongs to the rare class of type-Ibn SNe. This SN has a large number of Swift-XRT and Chandra observations. The SN is detected on multiple epochs and its X-ray light curve is presented in Figure 2 (see also Immler et al. 2008). The SN was detected in X-rays soon after maximum optical light and reached a maximum X-ray luminosity of about  erg s at  days after maximum optical light, times the SN rise time. SN 2006jc was observed by Chandra on several occasions. We reduced a 55 ks Chandra observation with 306 photons at the SN location taken 87 days after maximum light. Given the limited number of photons we did not attempt to fit complex models. We found that the spectrum is well fitted by an power-law and with negligible absorbing column density. The spectra and the best fit model are presented in Figure 3. Regardless of the exact spectral shape, the spectrum is hard. Marginalizing over all the free parameters, we find a 2- upper limit of  cm, in excess of the Galactic column density.

Figure 2.— Upper panel: The Swift (circles) and Chandra (squares) X-ray light curve extracted at the position of SN 2006jc. The unabsorbed flux was calculated using PIMMS assuming Galactic column density of  cm, and power-law spectrum. lower panel: The mean photon energy of the Swift-XRT X-ray observations in the 0.2–10 keV band as a function of time.

Figure 3.— spectrum of SN 2006jc, as observed at 87 days after maximum optical light, near or at peak X-ray luminosity. The top panel shows the binned X-ray data (points) with errors (1) along with the best-fit model, a photon index (90%-confidence) power-law, where , and with negligible line of sight absorption. The dashed line shows the constrained best-fit that has maximal (90%-confidence;  cm). The bottom panel shows the residuals to the best fit. The fit is acceptable, with for 12 degrees of freedom ().

SN 2011ht (Fig. 4): The SN took place about from the center of UGC 5460. It was observed on multiple epochs using Swift-XRT. Roming et al. (2012) reported a detection of an X-ray source at the SN position. The binned Swift-XRT X-ray light curve of this SN is shown in Figure 4. Apparently the light curve is rising, peaking at around 40 day after maximum optical light, and then declines. However, the uncertainties in the flux measurements are large and the light curve is consistent with being flat (i.e., a best fit flat model gives ). Moreover, recently Pooley (2012) reported on a 9.8 ks Chandra observation31 of this SN. He argued that the emission detected by Roming et al. (2012) which is shown in Figure 4 is from a nearby source found 4.7” from the SN location.

Figure 4.— The unabsorbed Swift-XRT light curve extracted at the position SN 2011ht in the 0.2–10 keV band, corrected for the aperture size and assuming  cm. The dashed gray line represents the 2- upper limit on the flux from two combined Swift-XRT observations obtained and days prior to maximum light.

SN 2010al (Fig. 5): This SN was found from the center of the spiral galaxy UGC 4286. The SN was observed in multiple epochs using Swift-XRT with a total integration time of 35 ks. It is detected in the combined image with a mean luminosity of  erg s. Figure 5 presents the binned Swift-XRT light curve of this SN. Although the light curve peaks around 30 days after maximum light, it is consistent with being flat (i.e., a best fit flat model gives ).

Figure 5.— The unabsorbed Swift-XRT light curve extracted at the position SN 2010al in the 0.2–10 keV band, corrected for the aperture size and assuming  cm.

SN 2005db: This SN was observed on three epochs, about two years after maximum light (676 to 695 days), using Swift-XRT. The combined image, with an exposure time of 13.6 ks, shows a faint source (5 photons) with a false alarm probability of per trial. However, given the fact that we have 305 observations, the false alarm probability over all trials is about . Using the Galactic column density (Table 1) and assuming a power-law spectrum , the unabsorbed flux is  erg s. Chandra observed this target twice at 722.7 and 1051.4 days after maximum light (PI: Pooley), with integrations of 3.0 and 5.0 ks, respectively. Using the same assumptions as above, we put a - upper limit on the unabsorbed flux of  erg s for both epochs. To conclude, given the uncertainties, the Chandra upper limits are consistent with the possible Swift detection. However, given that this source has a single detection we cannot firmly conclude that the detection is real.

PTF 11iqb: This SN has multi-epoch Swift-XRT observations taken between about and 28 days relative to maximum light. The SN is detected in a single epoch about 24 days after maximum light, with a false alarm probability of per trial. However, given that we are reporting 305 individual X-ray observations, the false alarm probability over all the trials is . The SN was not detected at the last epoch, 28 days after maximum light. Unfortunately, this object has no further Swift-XRT observations.

SN 2007pk (Fig. 6): This SN has a large number of Swift-XRT observations, as well as a Chandra observation (PI: Pooley). Immler et al. (2007) reported a tentative detection of this SN in the images taken between MJD 54417.09 and 54420.04. The light curve of the source extracted in the SN position is shown in Figure 6. The light curve shows a brightening, with a peak around MJD 54461, and a full width at half maximum of about 40 days. However, the SN is about from the center of the spiral host galaxy, NGC 579, and the centroids of the Swift-XRT positions in individual exposures are clustered around the galaxy nucleus, rather than the SN position. We note that emission from the center of NGC 579 is clearly detected in the Chandra observation and that there is some emission in the source position. However, this emission may be due to diffuse emission from NGC 579. Therefore, without conclusive evidence that the emission is from the SN, here we assume that the observed flare as well as the quiescent X-ray emission from the position of the source is due to AGN activity in NGC 579. In Table 1, we adopted an upper limit on the X-ray luminosity of SN 2007pk, which is based on the average luminosity observed from the direction of the source, presumably due to AGN activity.

Figure 6.— X-ray light curve, corrected for the aperture size, extracted at the position SN 2007pk. The emission, and flare, are likely due to AGN activity in the host galaxy NGC 579.

4. The model

Several recent works discuss the possibility of detecting X-ray emission from SN collisionless shocks in wind environments (Katz et al. 2011; Murase et al. 2011; Chevalier & Irwin 2012; Svirski et al. 2012). Here, we review the main processes relevant to shock breakout in wind-profile CSM (§4.1), emission from collisionless shocks, including the importance of bound-free absorption (§4.2), and the possibility of detecting radio emission (§4.3). In §5 we discuss our observations in the context of this model.

4.1. Shock breakout conditions in wind-profile CSM

Here, we assume that the CSM around the progenitor has wind-density profile , where is the mass-loading parameter, is the progenitor mass-loss rate, and is the progenitor wind speed. The Thomson optical depth, , due to an ionized progenitor wind between the observer and a spherical surface at radius from the star center is

(1)
(2)

Here is the opacity in units of 0.34 cm gr, is the mass-loss rate in units of 0.1 M yr, is the wind speed in units of 10 km s, and is the radius in units of  cm. We note that this relation is correct up to a correction factor of the order of unity (see Balberg & Loeb 2011). The photons in the radiation dominated and mediated shock from the SN explosion breakout when (e.g., Weaver 1976; Ofek et al. 2010), where is the speed of light and is the shock velocity. At this optical depth, the photon diffusion time becomes shorter than the hydrodynamical time scale (i.e., ) and the photons can diffuse outward faster than the ejecta. Therefore, the condition for the shock breakout to take place in a steady wind environment () is

(3)

where is the Thomson optical depth in units of 30, and is the radius in units of the solar radius.

Figure 7 shows the radius vs. the mass-loading parameter at which (Eq. 3; i.e.,  km s; black-solid line).

Figure 7.— Various regions in the radius vs. wind mass-loading parameter phase space, assuming . As indicated in the legend the black-solid line represent Thomson optical depth, (i.e., shock breakout with  km s). and are marked with the black dashed, and black dashed-dotted lines, respectively. The dotted black lines represent the integral of mass inside a given radius (i.e., ). The gray lines show constant bound-free absorption at a given energy. For reference, the radii of red supergiant (RSG; 500 R), the Sun and a massive white dwarf (WD; 0.005 R) are marked in circles on the line. See text for discussion.

For example, this plot suggests that the critical mass-loading, above which the shock will breakout in the wind environment, is  M yr v for a 500 R red supergiant. Assuming  km s, systems found above the solid line will have a shock breakout in a wind profile (the two cases discussed in Chevalier & Irwin 2011). Objects found in regions below the , black dashed-dotted line, will have a shock breakout within the stellar surface and the wind will not influence the diffusion of energy from the shock breakout. Finally, systems below the solid line () and above the dashed-dotted line () will have a shock breakout below the stellar surface, but the wind can play a role in the diffusion of the shock energy (e.g., Nakar & Sari 2010).

We note that, in order to form a type-IIn SN, we require optically thin material, which is ionized by the SN radiation field. Therefore, we speculate that type-IIn SNe can be found below and above the line, and that not all type-IIn SNe are powered by shock breakout in dense CSM environments.

4.2. X-ray emission from collisionless shocks

Katz et al. (2011) showed that if the shock width () at radius is of the order of (i.e., the Thomson optical depth vary on scales of the order of the radius ), the radiation mediated and dominated shock will transform into a collisionless shock, and hard X-ray photons will be generated. It is reasonable to assume that some of this energy will be produced in the 1–10 keV X-ray regime. We note that the exact spectrum was not calculated self-consistently and, therefore, is not known.

Chevalier & Irwin (2012) and Svirski et al. (2012) showed that, during the first several shock breakout time scales after the shock breakout, the optical depth is too large for the hard X-rays to escape. They found that the most efficient processes in blocking the hard X-rays ( keV) are likely: (i) cooling of the electrons by inverse Compton scattering on soft photons generated by the pre-shocked material; and (ii) Compton scattering. On average, for each Compton scattering, the photon loses a fraction of energy which is comparable to (e.g., Lang 1999). Here is the Boltzmann constant, is the electron temperature, and is the electron mass. Svirski et al. (2012) argued that when the Thomson optical depth declining to , the hard X-ray emission ( keV) will be Comptonized to the 1–10 keV X-ray band and may diffuse out of the CSM. Since the optical depth in a wind profile () decrease as (Equation 2), the collisionless shock signature will be observable in the X-ray regime only when the optical depth decreases by an order of magnitude (i.e., the shock radius increases by an order of magnitude). Given that the shock velocity falls as , they argued that this should happen roughly between 10 to 50 times the shock-breakout time scale. For reference, we show the line in Figure 7 (dashed line; assuming  km s). Such late-time emission may be relevant only if the collisionless shock is still important at late times and if the Comptonization remains the dominant process.

Chevalier & Irwin (2012) consider the effect of bound-free absorption. Since, above 0.5 keV, metals with a high ionization potential dominate the bound-free absorption, full ionization is required in order to avoid absorption by this process. They argued that an ionization parameter of is needed to achieve full ionization (including the metal atoms’ inner electrons). They estimate that full ionization is achieved for shocks with velocity  km s. In order to estimate the effect of bound-free absorption, we need to evaluate the density of the CSM.

Assuming a hydrogen-rich material, the particle density in a wind profile is given by

(4)
(5)

where is the mean number of protons per particle (mean molecular weight). For our order of magnitude calculation, we assume . In a wind profile, the column density between the radius and the observer is:

(6)

Assuming the gas in the pre-shocked wind is neutral, the bound-free optical depth in the 0.03 to 10 keV region is roughly given by (e.g., Behar et al. 2011)32:

(7)
(8)

where is the bound-free cross section as a function of energy , and is the energy in keV. This approximation is valid when the material is neutral. However, since above  keV, metals with a high ionization potential dominate the absorption, this formula is still valid, to an order of a magnitude, above 0.5 keV when some (or even one) of the inner electrons of the metals are bound. In Figure 7 we show the lines at which at 1 keV and 10 keV, and at which at 1 keV. Comparing equations 2 and 8 suggests that at the time of shock breakout the bound-free cross-section in the  keV range is larger than the Thomson cross section. We note that this may modify the properties of the shock breakout and its spectrum. Moreover, the line at 1 keV is located far below the line. This suggests that if the pre-shocked wind is not completely ionized (e.g.,  km s), then soft ( keV) X-ray emission is not expected in the simple case of a spherically symmetric wind () profile, even at late times. Moreover, the at 10 keV line is located slightly below the line. Therefore, bound-free absorption is likely important even at late stages. This may indicate that 10 keV X-rays may escape the wind on a time scale somewhat longer than suggested by Svirski et al. (2012), and that observations at energies above 10 keV may be more effective (e.g., by the NuStar mission; Harrison et al. 2010).

All the order of magnitude calculations presented so far assume a wind density profile with . However, we note that if the CSM profile falls faster than , or alternatively if the wind is not spherically symmetric or is clumpy, then the column density may fall faster than in some directions (e.g., Eq. 6), and it will enable the X-rays to escape earlier than predicted by Svirski et al. (2012) for the case.

4.3. Radio emission from collisionless shocks

The shock going through the CSM may generate synchrotron radiation peaking at radio frequencies (e.g., Slysh 1990; Chevalier & Fransson 1994; Chevalier 1998; Horesh et al. 2012; Krauss et al. 2012; Katz 2012). However, if the material is ionized or partially ionized then the free-free optical-depth may block this radiation. In order to test if radio signature is expected, we need to estimate the free-free optical depth (e.g., Chevalier 1981) which, for a ionized CSM with a wind profile, is given by (Lang 1999, Equation 1.223)

(9)
(10)

where is the frequency in units of 10 GHz. The free-free optical depth is so large that radio emission is not expected. However, if in some regions the CSM density profile falls significantly faster than then the free-free absorption may be low enough for radio emission (e.g., synchrotron) to escape the CSM. We predict that if the CSM is ionized, then due to the effect of free-free absorption, the synchrotron radio emission generated in the shocked CSM may have a relatively steep radio spectrum. Therefore, it is preferable to search for this emission at high frequencies and late times after maximum light. However, the existence of radio emission is likely very sensitive to the exact properties, like density profile, symmetry, and homogeneity of the CSM. We speculate that good candidates for radio emission will be SNe in which the wind filling factor is low, or asymmetric, or alternatively when the wind is ejected in a relatively short eruption, and therefore may have a steep density profile (see §5).

We note that van Dyk et al. (1996) presented a search for radio emission among ten type-IIn SNe. None of the SNe in their sample was detected in radio. However, these observations were conducted between 2 yr to 14 yr after the SN explosion.

5. The observations in context of the model

As shown in Table 2 and Figures 15, several SNe in our sample show X-ray emission in the 0.2–10 keV range. Some SNe are presumably detected (and maybe peaking) in X-rays near maximum optical light (i.e., SN 2011ht; SN 2010al). However, the X-ray luminosity of SN 2011ht and SN 2010al at maximum visible light is about and , respectively, of their visible light luminosity. These X-ray luminosities are higher than predicted by Svirski et al. (2012; i.e. X-rays of optical) for the CSM shock breakout case. Moreover, although the X-ray light curves are consistent with a non-variable luminosity, it is possible that they are peaking near maximum optical light. Therefore, we conclude that the optical light curves of SN 2011ht and SN 2010al, which are type-IIn SNe, are likely not powered by a shock breakout in CSM.

SN 2010jl (Chandra et al. 2012a) and SN 2006jc (Immler et al. 2007) are detected in X-rays and are peaking at late times and , respectively. Moreover, the X-ray luminosity at maximum visible light is about of the visible-light luminosity. This is roughly consistent with the predictions of Svisrski et al. (2012). Therefore, we conclude that these two SNe are likely powered by shock breakout in CSM. We discuss these SNe in detail in §5.1-5.2.

Two other SNe in our sample, PTF 11iqb and SN 2005db, have marginal X-ray detections. Therefore, we cannot make any firm conclusion regarding the reality and nature of this emission.

As indicated in the column in Table 1, the rest of the SNe in our sample, including all the SLSN-I events, do not have late time observations and/or deep enough limits in order to evaluate their nature. Our upper limits are mostly obtained at early times () after the SN explosion, or at very late times (). We note that all the observations of the hydrogen-poor luminous SNe were obtained at after optical maximum light.

We note that recently Chandra et al. (2012b) reported on X-ray and radio observations of another type-IIn event, SN 2006jd. This event is listed as a type-IIb SNe in the IAUC SN list33 and, therefore, was not included in our sample. We note that the X-ray observations of this event started about a year after the explosion, so there is no measurement of during optical maximum light.

We conclude that deeper X-ray observations over longer periods of time, since maximum optical light, are required in order to understand the nature of type-IIn and SLSN-I events. The current null detection of hydrogen-poor luminous SNe in X-rays cannot be used to reject the CSM-interaction model proposed by Quimby et al. (2011a).

5.1. SN 2010jl (PTF 10aaxf)

The X-ray luminosity of SN 2010jl near optical maximum light is about  erg s, which is of its -band luminosity. For the case of shock breakout in CSM, Svirski et al. (2012) predicted that near optical maximum light the X-ray luminosity will be about of the optical luminosity. However, it is possible that the bolometric optical luminosity is higher (e.g., due to metal blanketing). Moreover, we note that for , the amount of material above the shock is smaller and the X-rays will be less effected by absorption at early times.

The Swift X-ray flux is rising with time, and is consistent with a power-law of the form , where is the time since optical maximum light. Svirski et al. (2012) predicted that the hard radiation (i.e., X-ray) component of a collisionless shock will rise as with is between to when the ejecta are colliding with a wind mass that is comparable to its own mass, and with between to when the wind is less massive than the ejecta. However, our observations constrain only the  keV range, while the hard component can emit at energies up to  keV. Moreover, Svirski et al. (2012) assumed that the bound-free absorption can be neglected even at late times (i.e., the CSM is completely ionized).

Chandra et al. (2012a) reported on the analysis of the X-ray spectrum of SN 2010jl (including the proprietary data). They measured a column density at their latest epoch of  cm, which is about 1000 times larger than the Galactic column density in the direction of SN 2010jl. Such a large bound-free absorption is expected if the shock velocity is below  km s (i.e., the metals are not completely ionized; Chevalier & Irwin 2012) and the mass-loading parameter is as large as expected from a wind shock breakout (e.g., Eq. 6).

Moreover, the column density decreases by a factor of between 346 and 405 days after optical maximum light. Equation 6 predicts that between these dates, assuming , the bound-free absorption should decrease by only %. In order to explain such a fast decline in column density over an increase in time, we suggest that in some regions above the shock breakout region the CSM may have a steep density profile. If indeed the CSM have steep power-law index, we predict that the free-free absorption will be low enough and it may be possible to detect late-time radio emission from this SN (see §4.3). A complete analysis of all the available data, including the proprietary data, is required in order to understand the luminosity and spectral evolution of this SN, and to give more firmer predictions.

5.2. SN 2006jc

The X-ray light curve of SN 2006jc (Fig. 2; Immler et al. 2008) peaked around  days after the explosion. The visible-light rise time of SN 2006jc was shorter than about 15 days (Foley et al. 2007; Pastorello et al. 2008). Therefore, if this SN is powered by the diffusion of shock energy and if we can approximate the shock breakout time scale as the visible-light rise time, then we can deduce that the X-rays peaked . The emission line width in the spectra of SN 2006jc suggests that the shock velocity was likely below 4000 km s (Foley et al. 2007). If this estimate is correct, then it is possible that the pre-shocked CSM of this SN is partially ionized (Chevalier & Irwin 2012). This may explain the fact that the X-ray spectra of SN 2006jc become softer with time (Fig. 2; Immler et al. 2008). As the shock propagates through the CSM the column density and, therefore, the bound-free absorption, both decreases. This is consistent with the expectation that the bound-free process will dominate the absorption of soft X-rays, and therefore, X-rays at 1 keV will escape only after the column density drops below  cm.

Itagaki (2006) reported that a possible eruption, with an absolute magnitude of about , took place at the position of SN 2006jc about two years prior to the SN explosion (see also Pastorello et al. 2008). This pre-SN outburst of SN 2006jc may have put in place a dense shell that could provide the medium required for the formation of a collisionless shock. Moreover, due to the eruptive nature of the event, the outer edge of this shell may follows a density profile that falls faster than , or has a relatively sharp edge. Such a density profile is required for the X-rays to escape the shocked regions at times earlier than the 10–50 shock-breakout time scale suggested by Svirski et al. (2012).

6. Conclusions

We present a search for X-ray emission from 28 SNe in which it is possible that the shock breakout took place within a dense CSM. Most SNe have been observed with X-ray telescopes only around maximum optical light. The SNe in our sample that do have late time observations, were either detected also at early times or were observed serendipitously at very late times. In that respect, our first conclusion is that a search for X-ray emission both at early and late time from SNe is essential to constrain the properties of the CSM around their progenitors.

Our analysis suggest that some type-IIn/Ibn SNe, most notably SN 2006jc and SN 2010jl, have optical light curves that are likely powered by a shock breakout in CSM, while some other type-IIn SNe do not. However, for most of the SNe in our sample, including all the SLSN-I events, the observations are not conclusive. Specifically, the lack of X-ray detection of SLSN-I events, cannot rule out the interaction model suggested in Quimby et al. (2011a; see also Ginzburg & Balberg 2012). We conclude that deeper observations at later times are required in order to further test this model.

Given the limits found in this paper and our current understanding of these events it will be worthwhile to monitor type-IIn SNe (as well as other classes of potentially interacting SNe; e.g., type-IIL) with X-ray and radio instruments at time scales times the rise time of the SN. We argue that in some cases bound-free absorption will play an important role at early and late times. Therefore, observations with the soon to be launched Nuclear Spectroscopic Telescope Array (NuSTAR; Harrison et al. 2010) in the 6–80 keV band may be extremely useful to test the theory and to study the physics of these collisionless shocks. Moreover, in the cases where bound-free absorption is important (e.g.,  km s; Chevalier & Irwin 2012), the spectral X-ray evolution as a function of time can be use to probe the column density above the shock at any given time, and deduce the density profile outside the shocked regions. We also argue that in some cases, if the CSM has steep density profiles (e.g., SN 2010jl), it may be possible to detect radio emission.

Finally, we note that Katz et al. (2011) and Murase et al. (2011) predict that the collisionless shocks will generate TeV neutrinos. These particles will be able to escape when the collisionless shock begins. The detection of such neutrinos using Ice-Cube (Karle et al. 2003) will be a powerful tool to test this theory and explore the physics of collisionless shocks.

We thank Ehud Nakar, Boaz Katz and Nir Sapir for many discussions. This paper is based on observations obtained with the Samuel Oschin Telescope as part of the Palomar Transient Factory project, a scientific collaboration between the California Institute of Technology, Columbia University, Las Cumbres Observatory, the Lawrence Berkeley National Laboratory, the National Energy Research Scientific Computing Center, the University of Oxford, and the Weizmann Institute of Science. EOO is incumbent of the Arye Dissentshik career development chair and is grateful to support by a grant from the Israeli Ministry of Science. MS acknowledges support from the Royal Society. AC acknowledges support from LIGO, that was constructed by the California Institute of Technology and Massachusetts Institute of Technology with funding from the National Science Foundation and operates under cooperative agreement PHY-0757058. SRK and his group are partially supported by the NSF grant AST-0507734. AG acknowledges support by the Israeli, German-Israeli, and the US-Israel Binational Science Foundations, a Minerva grant, and the Lord Sieff of Brimpton fund. The National Energy Research Scientific Computing Center, which is supported by the Office of Science of the U.S. Department of Energy under Contract No. DE-AC02-05CH11231, provided staff, computational resources, and data storage for the PTF project. PEN acknowledges support from the US Department of Energy Scientific Discovery through Advanced Computing program under contract DE-FG02-06ER06-04. JSB’s work on PTF was supported by NSF/OIA award AST-0941742 (“Real-Time Classification of Massive Time-Series Data Streams”). LB is supported by the NSF under grants PHY 05-51164 and AST 07-07633. MMK acknowledges generous support from the Hubble Fellowship and Carnegie-Princeton Fellowship.

Footnotes

  1. affiliation: Benoziyo Center for Astrophysics, Weizmann Institute of Science, 76100 Rehovot, Israel.
  2. affiliation: Department of Astronomy & Astrophysics, Pennsylvania State University, University Park, PA 16802
  3. affiliation: Department of Astronomy, University of California, Berkeley, Berkeley, CA 94720-3411.
  4. affiliation: Department of Physics, University of Oxford, Denys Wilkinson Building, Keble Road, Oxford OX1 3RH, UK.
  5. affiliation: Racah Institute of Physics, The Hebrew University of Jerusalem 91904, Israel
  6. affiliation: National Radio Astronomy Observatory, P.O. Box O, Socorro, NM 87801
  7. affiliation: Division of Physics, Mathematics and Astronomy, California Institute of Technology, Pasadena, CA 91125.
  8. affiliation: LIGO laboratory, Division of Physics, California Institute of Technology, MS 100-36, Pasadena, CA 91125
  9. affiliation: Kavli IPMU, University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa-shi, Chiba, 277-8583, Japan
  10. affiliation: NASA-Goddard Space Flight Center, Greenbelt, Maryland 20771
  11. affiliation: Division of Physics, Mathematics and Astronomy, California Institute of Technology, Pasadena, CA 91125.
  12. affiliation: Benoziyo Center for Astrophysics, Weizmann Institute of Science, 76100 Rehovot, Israel.
  13. affiliation: Lawrence Berkeley National Laboratory, 1 Cyclotron Road, Berkeley, CA 94720.
  14. affiliation: Benoziyo Center for Astrophysics, Weizmann Institute of Science, 76100 Rehovot, Israel.
  15. affiliation: Department of Astronomy, University of California, Berkeley, Berkeley, CA 94720-3411.
  16. affiliation: Observatories of the Carnegie Institution for Science, 813 Santa Barbara St, Pasadena CA 91101 USA
  17. affiliation: Kavli Institute for Theoretical Physics, Kohn Hall, University of California, Santa Barbara, CA 93106.
  18. affiliation: Department of Physics, Broida Hall, University of California, Santa Barbara, CA 93106.
  19. affiliation: Department of Astronomy, University of California, Berkeley, Berkeley, CA 94720-3411.
  20. affiliation: School of Physics and Astronomy, Tel-Aviv University, Israel
  21. affiliation: Benoziyo Center for Astrophysics, Weizmann Institute of Science, 76100 Rehovot, Israel.
  22. affiliation: Spitzer Science Center, California Institute of Technology, M/S 314-6, Pasadena, CA 91125
  23. affiliation: Division of Physics, Mathematics and Astronomy, California Institute of Technology, Pasadena, CA 91125.
  24. affiliation: Division of Physics, Mathematics and Astronomy, California Institute of Technology, Pasadena, CA 91125.
  25. affiliation: Spitzer Science Center, California Institute of Technology, M/S 314-6, Pasadena, CA 91125
  26. slugcomment: Draft of March 21, 2018
  27. http://www.astro.caltech.edu/ptf/
  28. Weizmann Interactive Supernova (data) REPository; http://www.weizmann.ac.il/astrophysics/wiserep/
  29. Typically the background level is known very well.
  30. http://cxc.harvard.edu/toolkit/pimms.jsp
  31. At the time of the writing this paper, this observation was proprietary.
  32. This approximation deviates by a factor of two from a more accurate calculation (e.g., Morrison & McCammon 1983).
  33. http://www.cbat.eps.harvard.edu/lists/Supernovae.html

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