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The type Ia supernova SNLS-03D3bb from a super-Chandrasekhar-mass white dwarf star

Abstract

The accelerating expansion of the Universe, and the need for dark energy, were inferred from observations1,2 of type Ia supernovae. There is a consensus that type Ia supernovae are thermonuclear explosions that destroy carbon–oxygen white dwarf stars that have accreted matter from a companion star3, although the nature of this companion remains uncertain. These supernovae are thought to be reliable distance indicators because they have a standard amount of fuel and a uniform trigger: they are predicted to explode when the mass of the white dwarf nears the Chandrasekhar mass4 of 1.4 solar masses (M). Here we show that the high-redshift supernova SNLS-03D3bb has an exceptionally high luminosity and low kinetic energy that both imply a super-Chandrasekhar-mass progenitor. Super-Chandrasekhar-mass supernovae should occur preferentially in a young stellar population, so this may provide an explanation for the observed trend that overluminous type Ia supernovae occur only in ‘young’ environments5,6. As this supernova does not obey the relations that allow type Ia supernovae to be calibrated as standard candles, and as no counterparts have been found at low redshift, future cosmology studies will have to consider possible contamination from such events.

Main

The supernova SNLS-03D3bb (SN 2003fg) was discovered on 24 April 2003 ut as part of the Supernova Legacy Survey (SNLS). Its redshift is z = 0.2440 ± 0.0003, determined from host galaxy [O ii], [O iii], Hα, and Hβ emission lines. A finding chart and observational details can be found in the Supplementary Information. From the lightcurve (Fig. 1) we measure a peak magnitude in the rest frame V band, V = 20.50 ± 0.06 mag. This corresponds to an absolute magnitude of MV = -19.94 ± 0.06 (H0 = 70 km s-1 Mpc-l, ΩM = 0.3, flat Universe). SNLS-03D3bb falls completely outside the MV distribution of low-z type Ia supernovae (ref. 7), and is 0.87 mag (a factor of 2.2) brighter than the median. We note that neither changes in the Hubble constant nor ΩM significantly affect this brightness difference. Asphericity may account for variations in luminosity in a type Ia supernova at the 25% level, but not by a factor of two8,9. SNLS-03D3bb also does not follow the lightcurve width–luminosity relationship10 for type Ia supernovae that allows them to be calibrated as standard candles—it is too bright for its lightcurve width (‘stretch’, s = 1.13) by 0.61 ± 0.14 mag (4.4σ).

Figure 1: The lightcurve of SNLS-03D3bb.
figure1

We fitted k-corrected27 template lightcurves to the observed photometry of SNLS-03D3bb, then transformed the peak magnitudes back to the Johnson–Cousins28 BV magnitudes in the Vega system. We found peak magnitudes of B = 20.35 mag and V = 20.50 ± 0.06 mag from a simultaneous fit to g′ and r′ data. The error (s.d.) consists of 0.04 statistical error and 0.04 k-correction error. A lightcurve template was fitted using the stretch method1 (stretching the time axis of a template lightcurve by a stretch factor of s = 1.13). The epoch of maximum light relative to the rest-frame B band was determined from a simultaneous fit to all of the data. At maximum, we use only the V-band value to compare to other supernovae, because it is the best constrained. Data past +35 days were not used in the fit. The arrows are 3σ upper limits.

Type Ia supernovae are powered exclusively by the decay of 56Ni and its decay product 56Co (ref. 11), requiring 0.6M of 56Ni to reproduce a normal type Ia supernova12,13,14,15. Because SNLS-03D3bb is 2.2 times overluminous, we infer that it has 1.3M of 56Ni. Such a large 56Ni mass is not possible if the progenitor is limited to the Chandrasekhar mass. Even models that burn the entire 1.4M to nuclear statistical equilibrium via a pure detonation produce only 0.92M of 56Ni, with the remainder comprising other iron-peak elements16. At least 40% of the type Ia supernova must be elements other than 56Ni to reproduce observed spectra14,17; this implies a white-dwarf mass of 2.1M. Some authors find that rapid rotation may support such a massive white dwarf18. The merger of two massive white dwarfs could also produce a super-Chandrasekhar-mass product19,20.

This simple estimation of the nickel mass is supported by a more detailed calculation using the principle that the luminosity at maximum light is proportional to the instantaneous rate of radioactive decay21,22. The implied Ni mass is23,24: , where Lbol is the bolometric luminosity at maximum light (the luminosity integrated from the ultraviolet to the infrared), and α is the ratio of bolometric to radioactivity luminosities, near unity. is the radioactivity luminosity per solar mass of 56Ni from its decay to 56Co and subsequent decay to 56Fe: where tR is the time in days for the supernova to rise from explosion to maximum light. Using tR = s × 19.5 days (ref. 25), for SNLS-03D3bb, we obtain tR = 22 days (see Supplementary Information for the effect of a shorter rise). We use α = 1.2 as a conservative value, although for high 56Ni masses, α may be lower, because nickel above the photosphere will not contribute to the luminosity23.

To convert our V magnitude into a bolometric equivalent, we use a synthetic spectrum calculated to match the observed ultraviolet + optical spectrum (Fig. 3), but extended into the infrared (13% of the bolometric luminosity is from the infrared extrapolation). The bolometric correction (mbc) is 0.07 ± 0.03 mag, such that Mbol = MV + mbc = -19.87 ± 0.06 mag. Using these numbers, we calculate MNi = 1.29 ± 0.07M for SNLS-03D3bb, in agreement with the simple scaling argument used earlier. The quoted error is from the statistical, k-correction, and bolometric correction errors added in quadrature. SNLS-03D3bb has a significantly larger bolometric luminosity and implied 56Ni mass compared to low-redshift supernovae (Fig. 2).

Figure 3: Keck LRIS spectrum of SNLS-03D3bb at two days after maximum light compared to a spectrum of the normal type Ia supernova SN 1994D.
figure3

Also plotted is a SYNOW fit to the data, with dominant ions labelled. SYNOW is a parameterized resonance-scattering code, allowing the user to adjust optical depths, temperatures, and velocities to aid in the identification of supernova lines29. SYNOW parameters are listed in the Supplementary Information. SNLS-03D3bb shows the lines of intermediate-mass elements typically seen in a type Ia supernova at maximum light— Si ii, S ii, and Ca ii,—but in SNLS-03D3bb the velocity of the lines is lower than usual. The line at 415 nm appears to be C ii, but the other predicted carbon features cannot be clearly identified owing to the lower signal-to-noise ratio of the spectrum in the red. No other identification could be found for 415-nm feature.

Figure 2: Bolometric luminosity versus implied 56Ni mass for SNLS-03D3bb and low-redshift type Ia supernovae7.
figure2

The low-redshift type Ia supernovae were fitted using the same techniques as those used for SNLS-03D3bb: the bolometric luminosity was determined using the peak magnitude in the V band from a simultaneous fit to B- and V-band data. For the low-redshift type Ia supernovae, we integrated the s = 1 supernova template27 to obtain a bolometric correction of 0.06 mag and an uncertainty (s.d.) of 0.05 mag for the combined bolometric and k-correction error. The solid line represents a normal s = 1 type Ia supernova, with a rise time (tR) of -19.5 days, while dotted lines show s = 0.9 (tR = 17.6) and s = 1.1 (tR = 21.5). Low-luminosity type Ia supernovae have lower stretches, and thus shorter rise times, resulting in less 56Ni for a given luminosity, while high-stretch type Ia supernovae show opposite behaviour. The dashed line shows an upper limit for the expected 56Ni mass in a Chandrasekhar-mass type Ia supernova, obtained by burning the entire white dwarf to iron-peak elements in a detonation16.

SNLS-03D3bb also has an unusually low ejecta velocity, as shown in the Keck spectrum taken two days after maximum light (Fig. 3). With a Si ii velocity of 8,000 ± 500 km s-1, it falls well outside the range of velocities seen for this feature at maximum light (Fig. 4). This is hard to understand in the Chandrasekhar-mass model, which predicts higher velocities for more luminous type Ia supernovae, in contrast to the unusually low velocities in SNLS-03D3bb.

Figure 4: Inferred Ni mass versus Si ii 615 nm velocity.
figure4

a, The data at maximum light26. Ni masses are converted as described in the text using Mbol = MB + 0.2. Filled circles are from early-type galaxies (E or S0), while open circles are from late-type galaxies (Sa–Irr). Circle size is proportional to vSi/Si, where is the rate of change of the velocity of the Si ii feature. There is no measurement for SNLS-03D3bb. b, Kinetic-energy velocity of type Ia supernovae versus Ni mass for 1.4M models with different fractions of unburned carbon (fC). This unburned fraction should not be much higher than 20% because carbon is rarely seen in type Ia supernova spectra30. Overplotted symbols are vSi for low-redshift type Ia supernovae26 extrapolated to 40 days after maximum (correcting for stretch). For SNLS-03D3bb we use from its closest neighbour. The error bar reflects the range if an average value of is used. SNLS-03D3bb is not consistent with the 1.4M model. c, As above, but showing that Mwd ≈ 2M models can explain SNLS-03D3bb. Less extreme super-Chandrasekhar-mass models are consistent with the low-redshift data. The three low 56Ni supernovae are not necessarily of super-Chandrasekhar mass—their large values of Si make projections to 40 days uncertain.

The kinetic energy (EK) of a type Ia supernova arises from the difference between the nuclear energy (En) obtained from the synthesis of elements via fusion in the explosion and the binding energy (Eb) of the white dwarf13. Thus the kinetic energy velocity is: carbon–oxygen white dwarf is 0.5 × 1051 ergs (ref. 14). For a 2M white dwarf and a central density of 4 × 109 g cm-3, the binding energy is 1.3 × 1051 ergs (ref. 18).

There are only three classes of elements in a type Ia supernova (iron-peak elements, intermediate mass elements, and unburned carbon and oxygen), so a simple model can be developed for the nuclear-energy generation, En. Burning a mixture of equal parts carbon and oxygen to the iron peak produces EFe = 1.55 × 1051 erg M-1, while the synthesis of 28Si produces 76% as much energy13. Thus: En = EFeMwd(fFe + 0.76fIME), where Mwd is in solar masses, and fFe and fIME are the fractional compositions of iron peak and intermediate-mass elements. If fFe and fIME do not sum to one, the remainder is the fraction of unburned carbon and oxygen (fC), which does not contribute to the nuclear energy. The 56Ni makes up approximately 70% of iron-peak elements14,16, so we adopt MNi = 0.7MwdfFe, where MNi is the mass of 56Ni.

In the Chandrasekhar-mass model, more-luminous supernovae, with more 56Ni, have a higher vke (Fig. 4). Increasing the fraction of unburned carbon and oxygen, fC, can lower the kinetic energy, perhaps accounting for some of the dispersion in type Ia supernova velocities. However, this also lowers the available 56Ni, so it cannot account for the low velocity seen in SNLS-03D3bb.

The kinetic energy gives the velocity of the supernova averaged over the entire mass, approximately equivalent to the velocity a few weeks after maximum light13. The most appropriate observational signature of this velocity is unclear, because type Ia supernova line velocities change with time, and different ions can have different relative velocities. We find good agreement between the Si ii velocity at 40 days after maximum light26 and the theoretical kinetic-energy velocity, but we emphasize that this is an imperfect comparison.

A super-Chandrasekhar mass reproduces the low velocities seen in SNLS-03D3bb (Fig. 4). Because Chandrasekhar models with more Ni produce higher velocities, the low velocities of SNLS-03D3bb imply an increased progenitor binding energy and thus a larger total mass. As a caveat, we note that this simple calculation is only intended to illustrate general trends. Future theoretical studies will have to assess such complications using different ions, different white-dwarf density structures, and a wider range of binding energies.

Super-Chandrasekhar-mass type Ia supernovae should be more likely to occur in a young stellar population, where the most massive stars exist19,20. The low-mass, star-forming host of SNLS-03D3bb is consistent with this scenario (see Supplementary Information). Thus, the apparent existence of super-Chandrasekhar-mass type Ia supernovae may explain why the most luminous type Ia supernovae only occur in young stellar environments5,6. The standard Chandrasekhar-mass model offers no explanation for this behaviour, because the total amount of fuel and triggering mechanism should be independent of the mass of the progenitor stars.

Supernovae such as SNLS-03D3bb will have to be screened out in cosmological studies. Younger stellar environments produce more-luminous supernovae, so as the mean stellar age decreases with redshift the mean properties of type Ia supernovae will change5. This can be calibrated if all supernovae obey the same stretch–luminosity relationship, but SNLS-03D3bb does not. Its peculiarity was so obvious that it was excluded from the SNLS cosmological result7, but less extreme objects could lurk in supernova samples. Future cosmology studies will need to carefully scrutinise type Ia supernovae from young populations to see whether they obey the same lightcurve width–luminosity relationship as other type Ia supernovae.

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Acknowledgements

SNLS relies on observations with MegaCam, a joint project of CFHT and CEA/DAPNIA, at the Canada–France–Hawaii Telescope (CFHT). We used data products from the Canadian Astronomy Data Centre as part of the CFHT Legacy Survey. Some data were obtained at the W. M. Keck Observatory. We acknowledge support from NSERC, NERSC, CIAR, CNRS/IN2P3, CNRS/INSU, CEA and the DOE.

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Reprints and permissions information is available at www.nature.com/reprints. The authors declare no competing financial interests.

Correspondence to D. Andrew Howell.

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Supplementary Notes

This file contains Supplementary Methods, Supplementary Figure, Supplementary Table and additional references. (PDF 141 kb)

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