The Keck/SDSS SN Samplesdss2008.uchicago.edu/depot/foley-ryan.pdf · Fig. 8.Ñ (top panel):...

The Keck/SDSS SN Sample

Ryan FoleyUC Berkeley (CfA soon)

Alex Filippenko, Rick Kessler, Josh Frieman, SDSS-II SN Survey

– 20 –

! ! ! ! !

"#"

"#$

"#%

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"#'

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)*+,-./*!0 !

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Fig. 8.— (top panel): Composite spectrum created from our Keck/SDSS sample (blackcurve) compared to the maximum-light low-redshift composite spectrum from Foley et al.

(2007) (light-green curve). The spectra are scaled to match in the region 4500 ! ! ! 7500 A.The grey and dark-green regions are the 1" boot-strap sampling errors for the SDSS and low-redshift composite spectra, respectively. (second panel): The grey region is the 1" boot-strap

sampling region for the SDSS composite spectrum. The light-green curve is the residual ofthe SDSS and low-redshift composite spectra. The dark-green region is the residual of the

SDSS composite spectrum and the low-redshift 1" boot-strap sampling region. (third panel):The number of individual spectra contributing to each wavelength bin in the composite

spectra. (fourth panel): The average phase relative to maximum brightness as a functionof wavelength. (fifth panel): The average value of ! as a function of wavelength. (bottompanel): The average redshift as a function of wavelength.

1993ApJ...413L.105PSupernova Luminosity

Phillips 1993SNe Ia are not perfect standard candles

But the luminosity is calibrated by light-curve shape

– 57 –

Fig. 9.— Luminosity distance modulus vs. redshift for the ESSENCE, SNLS, and nearby

SNe Ia for SALT. For comparison the overplotted solid line and residuals are for a !CDM(w, "M, "!) = (−1, 0.27, 0.73) Universe.

Is SN Ia Evolution a Problem?

Wood-Vasey et al. 2007

SN Cosmology is a differential measurement

We assume that low-z and high-z SNe Ia follow the same luminosity vs. light curve shape relationship

If SNe Ia evolve, this relationship may break down leading to incorrect parameters

galactocentric distance, although perhaps the residual scatter isdecreased at large separations. There is also a hint that SNe Ia inelliptical hosts have slightly negative residuals after MLCS2k2correction (meaning that they are corrected to be too bright /toonearby); however, the weighted average residual is only!0:06 "0:06 mag, consistent with zero. Grouping together E and E/S0hosts (which show a similar paucity of slowly declining SNe Ia inthe top left panel) yields a weighted mean residual of !0:02 "0:04 mag. The observed scatter in the residuals of the early-typehosts is less than the overall sample, with ! # 0:11 mag for el-liptical hosts only and ! # 0:13 mag for E, E/S0, and S0 hosts; alarge fraction of this scatter may be from peculiar velocities, whichshould contribute$0.09mag to the dispersion for these galaxies.

Gallagher et al. (2005) have approached these issues in moredetail, with integrated spectra of the host galaxies of many of

these SNe Ia, allowing them to correlate SN properties (beforeand after MLCS2k2 correction) with additional parameters suchas host metallicity, star formation rate, and star formation history.Their results suggest no clear correlations with the MLCS2k2Hubble flow residuals, although there is marginal evidence for arelation between the residuals and host metallicity.

6.2. A Hubble Bubble?

Zehavi et al. (1998) presented evidence for a large local voidbased on SN Ia distances that suggested a monopole in the pe-culiar velocity field. They found that the Hubble constant esti-mated from SNe Ia within 70 h!1Mpc was 6:5% " 2:2% higherthan H0 measured from SNe Ia outside this region, assuming aflat !M # 1 universe. The significance of this void decreasesin the current concordance cosmology (!M # 0:3,!" # 0:7) to

Fig. 15.—Correlations with host galaxy morphology and projected galactocentric distance (GCD). The top panels show mean values of the light-curve shapeparameter# for the full sample, and the middle panels show the distribution of host galaxy extinction A0

V . The bottom panels show the residuals relative to the best-fitHubble line for the Hubble flow sample only (after MLCS2k2 correction for luminosity differences and extinction). Projected GCDs are calculated using the angularoffsets presented in Table 1, with angular diameter distances (!M # 0:3, !" # 0:7) calculated from the redshift for objects with czCMB % 2500 km s!1 and from theMLCS2k2 SN distances for objects with lower recession velocities.

JHA, RIESS, & KIRSHNER140 Vol. 659

galactocentric distance, although perhaps the residual scatter isdecreased at large separations. There is also a hint that SNe Ia inelliptical hosts have slightly negative residuals after MLCS2k2correction (meaning that they are corrected to be too bright /toonearby); however, the weighted average residual is only!0:06 "0:06 mag, consistent with zero. Grouping together E and E/S0hosts (which show a similar paucity of slowly declining SNe Ia inthe top left panel) yields a weighted mean residual of !0:02 "0:04 mag. The observed scatter in the residuals of the early-typehosts is less than the overall sample, with ! # 0:11 mag for el-liptical hosts only and ! # 0:13 mag for E, E/S0, and S0 hosts; alarge fraction of this scatter may be from peculiar velocities, whichshould contribute$0.09mag to the dispersion for these galaxies.




Zehavi et al. (1998) presented evidence for a large local voidbased on SN Ia distances that suggested a monopole in the pe-culiar velocity field. They found that the Hubble constant esti-mated from SNe Ia within 70 h!1Mpc was 6:5% " 2:2% higherthan H0 measured from SNe Ia outside this region, assuming aflat !M # 1 universe. The significance of this void decreasesin the current concordance cosmology (!M # 0:3,!" # 0:7) to


V . The bottom panels show the residuals relative to the best-fitHubble line for the Hubble flow sample only (after MLCS2k2 correction for luminosity differences and extinction). Projected GCDs are calculated using the angularoffsets presented in Table 1, with angular diameter distances (!M # 0:3, !" # 0:7) calculated from the redshift for objects with czCMB % 2500 km s!1 and from theMLCS2k2 SN distances for objects with lower recession velocities.


galactocentric distance, although perhaps the residual scatter isdecreased at large separations. There is also a hint that SNe Ia inelliptical hosts have slightly negative residuals after MLCS2k2correction (meaning that they are corrected to be too bright /toonearby); however, the weighted average residual is only!0:06 "0:06 mag, consistent with zero. Grouping together E and E/S0hosts (which show a similar paucity of slowly declining SNe Ia inthe top left panel) yields a weighted mean residual of !0:02 "0:04 mag. The observed scatter in the residuals of the early-typehosts is less than the overall sample, with ! ¼ 0:11 mag for el-liptical hosts only and ! ¼ 0:13 mag for E, E/S0, and S0 hosts; alarge fraction of this scatter may be from peculiar velocities, whichshould contribute$0.09mag to the dispersion for these galaxies.




Zehavi et al. (1998) presented evidence for a large local voidbased on SN Ia distances that suggested a monopole in the pe-culiar velocity field. They found that the Hubble constant esti-mated from SNe Ia within 70 h!1Mpc was 6:5% " 2:2% higherthan H0 measured from SNe Ia outside this region, assuming aflat !M ¼ 1 universe. The significance of this void decreasesin the current concordance cosmology (!M ¼ 0:3,!" ¼ 0:7) to


V . The bottom panels show the residuals relative to the best-fitHubble line for the Hubble flow sample only (after MLCS2k2 correction for luminosity differences and extinction). Projected GCDs are calculated using the angularoffsets presented in Table 1, with angular diameter distances (!M ¼ 0:3, !" ¼ 0:7) calculated from the redshift for objects with czCMB % 2500 km s!1 and from theMLCS2k2 SN distances for objects with lower recession velocities.


Jha, Riess, & Kirshner 2007

– 35 –

!0.2

0

0.2

0 10 20 30

Hubble

resid

ual [m

ag]

Age [Gyr]

Least!Squares fit

!0.4

!0.2

0

0.2

!0.5 0 0.5 1

Hubble

resid

ual [m

ag]

[M/H]

Least!Squares fit

TBT03 fit

0

1

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6

7

!0.02 !0.01 0 0.01

Counts

(x10

3)

!

No correlation ruledout at 56%

0

2

4

6

8

!0.2 0 0.2 0.4

Counts

(x10

3)

!

No correlation ruled

out at 98%

Fig. 9.— SN Ia Hubble residual vs. luminosity-weighted metallicity (top left) and age

(bottom left). A general trend is found suggesting that more metal rich populations produceintrinsically fainter SNe Ia. A least-squares fit finds to the data places a slope for the trendat 0.26. The dotted line represents the predicted trend based on the analytical models of

Timmes, Brown, & Truran (2003). The results of our statistical tests are presented in theadjacent plots. We rule out a no-correlation result for the HR vs. metallicity data at the

98% confidence level. A least-squares fit to the HR vs. age plot yields a nominal slope of!0.03 with negligible significance.

Gallagher et al. 2008

Physical Motivation

2 Howell

0

1

2

3

4

5

6

Num

ber

0.6 0.8 1.0 1.2 1.4Stretch

A component

passive hosts

0

5

10

15

20

Num

ber

B component

star forminghosts

Fig. 1.— Stretch distributions of SNe from each componentof the two component model. Top: Prompt (B-component) SNeIa. Bottom: Delayed (A-component) SNe Ia. The distributionswere derived from those in S06, as described in the text. Best-fitGaussians are shown. The A-component Gaussian is centered ats = 0.945, with ! = 0.077. The B-component Gaussian is centeredat 1.071 with ! = 0.063.

definition of M(t) as S06 — it is the mass turned intostars and does not include mass loss from supernovae.

The prompt and delayed SN Ia components have di!er-ent stretch distributions (S06). To determine the stretchof each SN, here we use the SNe from S06, though we fita new lightcurve template to the data using the SiFTOmethod (Conley et al. 2007a,b). Because stretches arealways defined relative to the s = 1 template lightcurve,stretch values should only compared within a publica-tion. However, the stretches derived here are approx-imately 4% larger than those in Astier et al. (2006),largely due to the use of a narrower s = 1 template.

All SNe from passive galaxies (i.e. those with no mea-surable star formation rate) were assigned to the A com-ponent. Star forming galaxies have SNe Ia from bothcomponents, so the A distribution from passive galaxieswas scaled by mass and subtracted from the distributionof SNe Ia from star forming galaxies, leaving the B dis-tribution (as in S06). The resulting distributions, andGaussian fits are shown in Figure 1. Note that to con-serve the total number of SNe one should add the SNesubtracted from the B distribution back to the A dis-tribution – this is unnecessary for our purposes becausethe relative heights of the gaussians are normalized as afunction of redshift in the next step.

To estimate the expected stretch evolution with red-shift, we take the observed A and B distributions andscale them to the predicted relative values with redshiftfrom Fig. 10 of S06. Increasing cosmic star formationwith redshift produces a larger fraction of SNe from theprompt component. Stellar mass as a function of redshiftis determined by integrating the star formation historyfrom the earliest times, so the total stellar mass, and thenumber of SNe from the A component, decreases withincreasing redshift. The net result is that in the A + Bmodel the mean stretch increases from 0.98 at z = 0 to

TABLE 1"2 and KS test: data and models

"2 KSz A+B MM A+B MM

0-0.1 0.81 0.63 15% 39%0.1-0.75 0.64 0.83 30% 28%0.75-1.5 0.60 0.84 52% 35%

Note. — Cols. 2-3: the "2 per degree of freedom between thedata and the predictions of the A+B and modified Mannucci (MM)models. Cols. 4-5: KS-test probability that the data is drawn fromeach model. Bins with zero counts were assigned an error of 1.15,possibly underestimating the "2. The KS test is also imperfectbecause probabilities were derived for a single Gaussian, not thesum of two Gaussians as used here. In both cases the two modelcomponents were fixed by the A and B numbers in S06.

1.04 at z = 1.5.One caveat is that in the A+B model there is no time

dependence for the A component. SNe Ia from 10 Gyrold progenitors are just as likely as SNe Ia from 3 Gyrold progenitors. If 10 Gyr-old SNe Ia are actually morerare, the A + B model will overpredict the number of A-component SNe at z = 0, as they result from stars formedduring the high star formation rate in the early universe(see discussion in S06). As an alternative to the A + Bmodel we tested the two component SN Ia delay time dis-tribution from Mannucci, Della Valle, & Panagia (2006),which has an exponential decrease in supernova prob-ability from the delayed component with time. Thedrawback of this model is that the probability distri-bution is somewhat arbitrary. Also, rather than the50-50 split between prompt and delayed SNe chosen byMannucci, Della Valle, & Panagia (2006), here we scaleeach component by the A and B values measured by S06.This gives similar results to the A+B model, predictinga shift in mean stretch from 0.98 to 1.02 from z = 0!1.5.

3. COMPARISON TO OBSERVATIONS

In Figure 2 we compare the predicted stretch distribu-tions from the A+B model to the observed stretch distri-butions in three redshift bins from the low redshift dataused by Astier et al. (2006), the SNLS data in S06, andthe data of the higher-z supernova search (Riess et al.2007). All lightcurves have been refit here using the samemethod. We also tested the data against the modifiedMannucci, Della Valle, & Panagia (2006) model, but wedid not find it to be a better predictor of SN evolutionwith redshift (Table 1).

Each survey has di!erent selection e!ects – the mostserious for the current study is Malmquist bias, the ten-dency to discover only the brightest members of a groupnear the detection limit of a magnitude-limited survey.To minimize the e!ect, for each of the high redshiftsearches we only consider supernovae from a reducedvolume so that none of the supernovae used are nearthe magnitude limit. The SNLS regularly discovers SNeIa out to z > 1, but here we use only the subset withz " 0.75, where Malmquist bias is minimal (Astier et al.2006). Similarly, we only use Riess et al. (2007) SNe withz < 1.5, where the authors report their sample is com-plete (Strolger et al. 2004). Lowering the redshift cutto!to z = 1.2 does not change the average stretch for thehighest-z sample, but it reduces the sample size from 20to 13, and thus decreases the significance of the results.

Evolution in average SN Ia properties with redshift 3

0

2

4

6

8

!0.4 !0.2 0.0 0.2 0.4

0.0 < z < 0.1

magnitude difference

PromptDelayed

0

1

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5

6

0.7 0.8 0.9 1.0 1.1 1.2 1.3stretch

0.75 < z < 1.5

0

5

10

15

20

Nu

mb

er

0.1 < z < 0.75

Fig. 2.— Actual stretch distributions compared to predictionsfrom the A + B model. In each case the A + B model evaluatedat the median redshift of the distribution is shown. Top: SNe Iafrom z = 0!0.1 (median z = 0.026; N=50). Middle: SNLS SNe Iafrom z = 0.1!0.75 (median z = 0.55; N=99). Bottom: Riess et al.(2007) SNe Ia from z = 0.75!1.5 (median z = 1.12; N=20), with serrors < 0.2 (requiring s error " 0.1 reduces the sample to 16 andincreases the average stretch from 1.06 to 1.07). The vertical linegives the mean stretch for each distribution. The top axis convertsa stretch di!erence into a magnitude di!erence using ! = 1.5.

As an additional protection against selection bias, weonly consider SNe with s ! 0.7. SNe Ia with s < 0.7are both dim and spectoscopically peculiar, like SN1991bg (Filippenko et al. 1992), and have not yet beendetected at z > 0.2, probably because of a combination ofMalmquist bias and spectroscopic selection bias (Howell2001; Howell et al. 2005) – as redshift increases, and theangular size of the host galaxy decreases, and it becomesmore and more di!cult to spectroscopically identify suchfaint supernovae when blended with their bright, oftenelliptical, hosts. This cut removes 3 SNe Ia from the low-z sample [other low-z SNe are already removed becausewe only consider Hubble-flow SNe Ia, with z > 0.015, tobe consistent with Astier et al. (2006)].

In all cases we use only SNe Ia with at least 4 lightcurvepoints, and at least one detection before 10 rest-framedays after maximum light in the B-band, so that stretchis accurately measured.

Figure 2 shows that the average observed stretch in-creases with redshift, from s = 0.98 ± 0.02 at a medianredshift of z = 0.03, to s = 1.02 ± 0.01 at z = 0.56, ands = 1.06 ± 0.02 at z = 1.12. Simultaneously the per-centage of SNe Ia with s < 0.9 decreases from 24% to15% to 1.4%. The KS test gives a 2% probability thatthe lowest and highest redshift sample are drawn fromthe same distribution. The predicted distributions fromthe A + B model are overplotted. The observed trendsmatch the predictions of the empirically-based models —with increasing redshift fewer low-stretch SNe Ia are ob-

0.0 0.1 0.2 0.3 0.4 0.5 0.6!m

!2.0

!1.5

!1.0

!0.5

w

BAOSNLS yr 1yr 1 s split

Fig. 3.— Cosmological fits done with the Astier et al. (2006)sample (dashed), and with SNe Ia drawn from the same sample,but using only s < 1 SNe Ia at z < 0.4 and s # 1 SNe Ia at z #0.4 (solid). Combining with Baryon Acoustic Oscillation results(Eisenstein et al. 2005) produces the thick lines. The results areconsistent at the one sigma level.

served, and the mean SN Ia stretch increases. We findthe same result when this analysis is repeated with theSALT (Guy et al. 2005) and SALT2 (Guy et al. 2007)lightcurve fitters. These results are also consistent withthe findings of Astier et al. (2006), that the low-z samplehad an average stretch 97% that of SNLS SNe.

Astier et al. (2006) estimate distances from SNe Ia us-ing

µB = m!

B " M + !(s " 1) " "c,

where µB is the distance modulus, m!

Bis the peak B-

band magnitude, c is a color, and M , ! and " are param-eters fit by minimizing residuals on the Hubble diagram.Astier et al. (2006) found ! = 1.52, so a drift in averagestretch of 0.08 ± 0.026 from z = 0.03 to z = 1.12 resultsin a 12% drift in average intrinsic SN Ia luminosity overthis redshift range.

4. EFFECT ON COSMOLOGICAL STUDIES

Evolution in the SN population will not necessarily biascosmological studies, since SNe are only used in this wayafter correction for lightcurve shape. However, we can nolonger assume that any deficiencies in lightcurve widthcorrection schemes will average out under the assumptionthe distribution of SNe is similar over all redshifts. Ifthere is a systematic residual between low stretch andhigh stretch SNe when they are stretch corrected, thiscould cause a bias in the determination of cosmologicalparameters as the population evolves.

To test an extreme case of evolution, we fit the equa-tion of state parameter for dark energy, w and the mat-ter density, "M using the data from Astier et al. (2006),(dashed lines in Fig. 3) and compared it to a fit using thesame data, but retaining only s < 1 SNe Ia at z < 0.4and only s ! 1 SNe Ia at z ! 0.4 (solid lines). Thestrongly evolving subset gives estimates for w and "M

consistent with the full set, although the errors are largerbecause there are fewer SNe in the subset.

One worry with an evolving population is that stretch-magnitude or color-magnitude relations may evolve, i.e.! or " derived at one redshift may not be appropriate atanother. The values derived here for the strong evolutionsubset and the full set are consistent, but again a strongtest awaits a larger data set.

Measuring changes in w with time will require much

Howell et al. 2007LLL L

L

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Galaxy demographics change with redshift and SN properties change with galaxy type...Evolution may change the luminosity

Theoretical Motivationar

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Mon. Not. R. Astron. Soc. 000, 1 – 10 (2006) Printed 15 August 2006 (MN LATEX style file v2.2)

Cosmological Implications of the Second Parameter ofType Ia Supernovae

Philipp Podsiadlowski1!, Paolo A. Mazzali2,3, Pierre Lesa!re1,4,5, Christian Wolf1

and Francisco Forster11Dept. of Physics, Oxford University, Oxford, OX1 3RH, UK2Max-Planck Institut fur Astrophysik, Karl-Schwarzschildstr. 1, 85748 Garching, Germany3Istituto Naz. di Astrofisica-Oss. Astron., Via Tiepolo, 11, 34131 Trieste, Italy4Institute of Astronomy, Cambridge CB3 0HA, UK5LRA, 24 rue Lhomond, 75231 PARIS Cedex 05, France

15 August 2006

ABSTRACT

Theoretical models predict that the initial metallicity of the progenitor of a Type Iasupernova (SN Ia) a!ects the peak of the supernova light curve. This can cause adeviation from the standard light curve calibration employed when using SNe Ia asstandardizable distance candles and, if there is a systematic evolution of the metallic-ity of SN Ia progenitors, could a!ect the determination of cosmological parameters.Here we show that this metallicity e!ect can be substantially larger than has beenestimated previously, when the neutronisation in the immediate pre-explosion phasein the CO white dwarf is taken into account, and quantitatively assess the importanceof metallicity evolution for determining cosmological parameters. We show that, inprinciple, a moderate and plausible amount of metallicity evolution could mimic a "-dominated, flat Universe in an open, "-free Universe. However, the e!ect of metallicityevolution appears not large enough to explain the high-z SN Ia data in a flat Universe,for which there is strong independent evidence, without a cosmological constant. Wealso estimate the systematic uncertainties introduced by metallicity evolution in a "-dominated, flat Universe. We find that metallicity evolution may limit the precisionwith which #m and w can be measured and that it will be di$cult to distinguishevolution of the equation of state of dark energy from metallicity evolution, at leastfrom SN Ia data alone.

Key words: cosmological parameters – distance scale – supernovae: general – super-novae: Type Ia – galaxies: evolution

1 INTRODUCTION

The use of Type Ia supernovae (SNe Ia) as standardizablecosmological distance candles (Riess et al. 1998; Perlmutteret al. 1999; Tonry et al. 2003; Riess et al. 2004; Astier etal. 2006) relies on the empirical fact that there is a tightcorrelation between the supernova peak brightness and thewidth of the supernova light curve (Phillips 1993), i.e. thefact that, to lowest order, SN Ia light curves form a one-parameter family of curves. The driving parameter that de-termines the relation, assuming that the mass of the progen-itor white dwarf at explosion is constant, has been shownto be the opacity in the ejecta (Khokhlov, Muller & Hoflich1993; Hoflich et al. 1996, Mazzali et al. 2001). This is closelyrelated to the quantity of radioactive 56Ni synthesised in the

! E-mail: [email protected]

explosion, which is responsible for the SN luminosity. In re-cent years it has become apparent from a larger sample ofobserved supernovae that not all SN Ia light curves fit intothis one-parameter family, producing an intrinsic, scatter inthe Phillips relation (see the discussion in Mazzali & Podsi-adlowski [2006] and Benetti et al. [2004]). This immediatelyimplies that there must be more than one parameter con-trolling SN Ia light curves. The physical property in theprogenitor that determines the dominating (first) parame-ter still has not been clearly identified. On the other hand,from a theoretical point of view it seems unavoidable thatthe metallicity of the original supernova progenitor mustat least in part be responsible for the intrinsic scatter aboutthe mean Phillips relation. Timmes, Brown & Truran (2003)showed, using straightforward and uncontroversial nuclearphysics arguments, that the neutron excess in the immediateprogenitor white dwarf, which is a direct function of the ini-

arX

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6Mon. Not. R. Astron. Soc. 000, 1 – 10 (2006) Printed 15 August 2006 (MN LATEX style file v2.2)

Cosmological Implications of the Second Parameter ofType Ia Supernovae

Philipp Podsiadlowski1!, Paolo A. Mazzali2,3, Pierre Lesa!re1,4,5, Christian Wolf1

and Francisco Forster11Dept. of Physics, Oxford University, Oxford, OX1 3RH, UK2Max-Planck Institut fur Astrophysik, Karl-Schwarzschildstr. 1, 85748 Garching, Germany3Istituto Naz. di Astrofisica-Oss. Astron., Via Tiepolo, 11, 34131 Trieste, Italy4Institute of Astronomy, Cambridge CB3 0HA, UK5LRA, 24 rue Lhomond, 75231 PARIS Cedex 05, France

15 August 2006

ABSTRACT

Theoretical models predict that the initial metallicity of the progenitor of a Type Iasupernova (SN Ia) a!ects the peak of the supernova light curve. This can cause adeviation from the standard light curve calibration employed when using SNe Ia asstandardizable distance candles and, if there is a systematic evolution of the metallic-ity of SN Ia progenitors, could a!ect the determination of cosmological parameters.Here we show that this metallicity e!ect can be substantially larger than has beenestimated previously, when the neutronisation in the immediate pre-explosion phasein the CO white dwarf is taken into account, and quantitatively assess the importanceof metallicity evolution for determining cosmological parameters. We show that, inprinciple, a moderate and plausible amount of metallicity evolution could mimic a "-dominated, flat Universe in an open, "-free Universe. However, the e!ect of metallicityevolution appears not large enough to explain the high-z SN Ia data in a flat Universe,for which there is strong independent evidence, without a cosmological constant. Wealso estimate the systematic uncertainties introduced by metallicity evolution in a "-dominated, flat Universe. We find that metallicity evolution may limit the precisionwith which #m and w can be measured and that it will be di$cult to distinguishevolution of the equation of state of dark energy from metallicity evolution, at leastfrom SN Ia data alone.

Key words: cosmological parameters – distance scale – supernovae: general – super-novae: Type Ia – galaxies: evolution

1 INTRODUCTION

The use of Type Ia supernovae (SNe Ia) as standardizablecosmological distance candles (Riess et al. 1998; Perlmutteret al. 1999; Tonry et al. 2003; Riess et al. 2004; Astier etal. 2006) relies on the empirical fact that there is a tightcorrelation between the supernova peak brightness and thewidth of the supernova light curve (Phillips 1993), i.e. thefact that, to lowest order, SN Ia light curves form a one-parameter family of curves. The driving parameter that de-termines the relation, assuming that the mass of the progen-itor white dwarf at explosion is constant, has been shownto be the opacity in the ejecta (Khokhlov, Muller & Hoflich1993; Hoflich et al. 1996, Mazzali et al. 2001). This is closelyrelated to the quantity of radioactive 56Ni synthesised in the

! E-mail: [email protected]

explosion, which is responsible for the SN luminosity. In re-cent years it has become apparent from a larger sample ofobserved supernovae that not all SN Ia light curves fit intothis one-parameter family, producing an intrinsic, scatter inthe Phillips relation (see the discussion in Mazzali & Podsi-adlowski [2006] and Benetti et al. [2004]). This immediatelyimplies that there must be more than one parameter con-trolling SN Ia light curves. The physical property in theprogenitor that determines the dominating (first) parame-ter still has not been clearly identified. On the other hand,from a theoretical point of view it seems unavoidable thatthe metallicity of the original supernova progenitor mustat least in part be responsible for the intrinsic scatter aboutthe mean Phillips relation. Timmes, Brown & Truran (2003)showed, using straightforward and uncontroversial nuclearphysics arguments, that the neutron excess in the immediateprogenitor white dwarf, which is a direct function of the ini-

Additional MotivationREPORT OF THE

DARK ENERGY TASK FORCE

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III. We recommend that the dark energy program include a

combination of techniques from one or more Stage III projects

designed to achieve, in combination, at least a factor of three gain over

Stage II in the DETF figure of merit, based on critical appraisals of

likely statistical and systematic uncertainties.

!

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IV. We recommend that the dark energy program include a

combination of techniques from one or more Stage IV projects

designed to achieve, in combination, at least a factor of ten gain over

Stage II in the DETF figure of merit, based on critical appraisals of

likely statistical and systematic uncertainties. Because JDEM, LST,

and SKA all offer promising avenues to greatly improved

understanding of dark energy, we recommend continued research and

development investments to optimize the programs and to address

remaining technical questions and systematic-error risks.

!

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V. We recommend that high priority for near-term funding should be

given as well to projects that will improve our understanding of the

dominant systematic effects in dark energy measurements and,

wherever possible, reduce them, even if they do not immediately

increase the DETF figure of merit.

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968 LENTZ ET AL. Vol. 530

explosion is the same as in Nugent et al. (1997). For thetwo-model !ts at days 10 (which we compare to day [9with respect to B maximum) and 15 (compared to day [4with respect to B maximum) for W7 we have used lumi-nosity parameters of g \ 0.4 and g \ 0.8, respectively. Forthese two models the !ts are generally good, but with all ofthe premaximum spectra, the red edge of the Si II feature at6150 does not extend far enough to the red. For modelA!day 7 (!t to day [12) we chose the model with luminosityg \ 0.1. This model !ts the luminosity and most featureswell ; however, like the day 10 models, the Ca H]K featuredoes not extend blueward enough when compared to theearliest observations. Hatano et al. (1999) have modeled thesame spectra and found that the Ca H]K feature requiresionized calcium with velocities up to D40,000 km s~1. Wehave conducted experiments extending the density pro!le ofW7. This did not provide enough optical depth in the CaH]K feature to a†ect the line pro!le. Neither do ourmodels reproduce the secondary feature at 4700 thatA!Hatano et al. (1999) ascribe to Fe II absorption in the C]Olayer. For our postmaximum model, day 35 (15 days aftermaximum), we !t the observation with a model with lumi-nosity g \ 1.5. This model !ts, generally ; however, the largefeature at 4800 (probably Fe II) is not seen in the obser-A!vations. The observed Na I D line feature is missing, prob-ably as a result of de!ciencies in the sodium model atom,which has already been improved in the latest version ofPHOENIX. Here, we are interested in small di†erentiale†ects, so it was essential that all models be run with thecode version frozen.

3. METALLICITY OF UNBURNED C]O LAYER

The simplest e†ect of metallicity on the spectral forma-tion in SNe Ia is the change in the metal content of theunburned C]O layer neglecting changes in the densitystructure and deeper layers. In this section we examine thesee†ects independently of other e†ects of progenitor metal-licity on SNe Ia. Since ongoing supernova searches areexpected to discover SNe early, and early spectra probe theoutermost layers only, this approach is sensible and yieldsphysical insight that is somewhat model independent.

When we look at the overall UVOIR synthetic spectra(Figs. 1, 6, 8, 11, and 13) of the models with variations in theC]O layer metallicity, we see two consistent and signi!-cant e†ects : shifts in the UV pseudocontinuum level(expanded view for day 7 in Fig. 2) and variations in the Si II

line at 6150 (expanded views in Fig. 5).A!The general e†ect in the UV is the increase in the UV

pseudocontinuum level with decreasing metallicity. Simul-taneous is the redward (blueward) shift of most UV featureswith decreasing (increasing) metallicity. In the UV, the line-forming region is in the C]O layer. As the metallicitydecreases, the line-forming region must reach deeper intothe atmosphere to have the same line opacity, resulting insmaller line velocities. Modi!cation of the C]O layermetal abundance gives a classic surface cooling e†ect : lowertemperatures for higher metallicity. The higher tem-peratures of the lower metallicity C]O atmospheres givehigher thermal Ñuxes, moving the UV pseudocontinuumhigher with lower metallicity. The surface cooling and theresulting shifts in UV pseudocontinuum are evident at everyepoch. There is the complementary e†ect of additionalmetals increasing the line blocking. We make no attempt toseparate these two e†ects in this paper.

FIG. 1.ÈModels with various metallicities in C]O layer at 7 days afterexplosion. T hick solid line : 10 times the normal C]O metallicity ; thickdotted line : 3 times ; short-dashed line : normal ; long-dashed line : 1/3 ; dot-dashed line : 1/10 ; and thin solid line : 1/30.

The Si II line at 6150 (Fig. 5) shifts blueward withA!increases in metallicity for epochs through day 20. Theseshifts demonstrate that some line formation in this featuretakes place in the C]O layer. The earlier epochs show largevariations in the total depth of the feature, which impliesthat the line forms less in the incomplete burning zone withits large, unchanging silicon abundance and more in theC]O layer where the silicon abundance changes. At laterepochs these conditions are reversed, resulting in smallerchanges with C]O layer metallicity variation. These e†ectsare discussed further in ° 5.1.

The Mg II ““ h]k ÏÏ feature at 2600 (Fig. 2) does notA!move to the blue or red as the metallicity varies. Thedecrease in Mg II h]k feature strength with increasingmetallicity is caused by the increasing UV line blanketingfrom background line opacity in the C]O layer. The Mg II

absorption occurs mostly in the deeper, partially burnedlayer that is highly enriched in magnesium. We have con-!rmed this hypothesis by calculating diagnostic outputspectra without using any background opacity (see Baron et

FIG. 2.ÈExpansion of ultraviolet region of Fig. 1

Lentz et al. 2000

Metallicity doesn’t affect optical spectrumIt does affect the UV

But models differ on the direction

6 D. N. Sauer et al.

Figure 4. Same as Fig. 3 except that in this series the mass fractions of Ni,Co, and Fe are increased together above 13 000 km s!1. Here the modelswith lower metals lead to an increased flux in wavelengths below "2500Åwhile in the part of the spectrum red of that wavelength the models withincreased Fe-group elements show more flux.

This model shows a clear decrease of the UV flux with an increaseof the metallicity.

In Fig. 6 the optical part of the model spectra are shown (forclarity we plot only every second model). It can be seen that theoptical part of the spectrum is less a!ected by the change in metal-licity than the UV.

To get an estimate for the magnitude of this e!ect, we com-pared the variation of the integrated flux of di!erent models. Forthis exercise we used the passband filters of the WFPC2 instru-ment, which are shown in the upper panel of Fig. 3 and Fig. 4.Fig. 7 shows the di!erence in magnitudes for the various filters rel-ative to the base model. In the model series shown in the upperpanel, a higher Fe abundance corresponds to more flux in the UVband-passes. In the most extreme case, the di!erence amounts to0.7mag. If all Fe-group elements are varied together (lower panel)the variation for increasing metal abundance is less pronounced.Instead, the models with lower metal abundance show a strong in-crease in flux in some filters. Generally this shows that the varia-tion of the flux with an increase of metal content can go both ways,brighter and dimmer, depending on the wavelength covered by therespective filter.

5 DISCUSSION

In this section we discuss the physical mechanisms that form theUV flux in SN Ia spectra and lead to the relationships between theobserved flux and the physical conditions in the ejecta discussed inthe previous section.

Figure 5. This model series corresponds to the one shown in Fig. 4 butat an earlier epoch of t = 15 d with L = 8.55 # 1042 erg s!1 and vph =10 000 km s!1.

Figure 6. The optical spectrum of the model series of Fig. 3 and Fig. 4. Forclarity only every second spectrum is shown here. The optical spectrum issignificantly less a!ected by the change in metal content than the UV part.

c$ 0000 RAS, MNRAS 000, 000–000

FIG. 7.ÈComparison of light curves in B (left) and V (right) of the delayed detonation models DD21c, DD23c, and DD24c with otherwise identicalparameters but di†erent C/O ratios and metallicities relative to solar (C/O; of (1 ; 1), 1), and (1 ; The di†erences relative to DD21c and theR

Z) (23 ; 13).

monochromatic light curves are given in the upper and lower plots, respectively.

FIG. 8.ÈComparison of synthetic NLTE spectra at maximum light for initial compositions of solar and solar (1o3), respectively, for model DD21c13(upper graph). The standard Johnson !lter functions for UBV and R are also shown. In the lower graph, the radius as a function of wavelength where themonochromatic optical depth reaches 0.1 and 1, respectively, is given for DD21c.

Höflich et al. 1998

Sauer et al. 2008

little evidence of evolution in the mean spectroscopic propertiesover this interval, confirming, with higher precision, earlier sug-gestions based largely on lower S/N data. The primary limitationin this aspect of our paper is the paucity of high-quality UV dataat low redshift.

Most of the earlier constraints on the possible evolution ofSNe Ia were largely based on consideration of photometric mea-sures. For example, Riess et al. (2004) derived the rest-frameU ! B and B! V distributions of a heterogeneous sample ofSNe Ia at maximum light over z < 1:5 as a proxy for the moreprecise constraints possible with spectroscopic data. They ruled outany color evolution greater than 0.02 mag in the mean in U ! B.

Spectroscopic data are substantially more precise in trackingevolution and dispersion than broadband photometry, for tworeasons. Foremost, detailed differences in key line diagnosticsare lost in photometric data. Secondly, no assumptions need bemade about k-corrections in gathering and comparing data over abroad redshift range.

At the time this Keck survey was underway, there was nocomprehensive spectroscopic data set against which our UV datacould be compared, other than the mean template published byNugent et al. (2002) and discussed in x 4.3. However, very re-cently, as a result of a survey of high-redshift SNe Ia with theACS grism on board HST, Riess et al. (2007) have published themean rest-frame UV spectrum of 13 SNe Ia more distant (z > 1)than those studied here; the average redshift of this sample isz " 1:3. In this study, used to demonstrate the presence of darkenergy at z > 1, it is claimed that the sample-averaged spectralenergy distributions observed at z >1 are consistent with thatobserved locally, and that any spectral evolution is still undetected.However, no quantitative statement is made to support this claim.

We believe that the spectral comparison undertaken by Riesset al. (2007) is less precise than that undertaken in this study, andhence less valuable as a means of justifying the continued use oflocal relations in constructing SN Ia Hubble diagrams. Foremost,the ACS grism data samples SNe Ia over a much wider range ofphase than the comparisons undertaken here. Although the in-dividual spectra comprising the mean z " 1:3 ACS spectrum arenot tabulated by Riess et al. (2007), their phases range over atleast 15Y20 days and almost all are postYmaximum light, com-pared with the narrow phase rangemaximum-light (!4 to +4 days)and preYmaximum light (<!4 days) comparisons presented here.A second consideration is the inevitable poor S/N of these veryhigh redshift spectra which precludes detailed comparisons.

Aswe have shown in Figures 5 and 6, the local template (Nugentet al. 2002) is a poor basis from which to make such comparisons.More representative data sets are needed for reliable claims. In-ternally within our own data, Figure 7 reveals little evolutionalthough the redshift baseline is small. It will be important to ex-amine the case for evolution at 0:2 < z < 1:5 by combining in aconsistent manner both the Riess et al. (2007) data set with thatpresented here, as well as with higher redshift SNLS SNe Ia (outto z # 1) observed during the routine survey spectroscopic screen-ing (e.g., Bronder et al. 2007).

The second significant finding in this study is the marked in-crease in the dispersion among our SNe Ia shortward of 3300 8(Fig. 10). Although the scatter inU is comparable to that seen inlocal data, it increases significantly at shorter wavelengths. Evenallowing for differing amounts of dust extinction within eachhost galaxy, this amounts to more than a factor of 2 variation incontinuum flux at maximum light. Although theoretical modelspredict a strong sensitivity to metallicity variations at this wave-length, our variations are also considerably larger than thosepredicted. We have demonstrated that a color correction based

on the B! V color of the SN and either a CCM Milky Way ex-tinction law or a SALTcolor law can only marginally reduce, andnot eliminate, this UV dispersion.We have found it hard to isolate the physical causes of this

significant UV dispersion.We confirm earlier work at longer wave-lengths (k ’ 3500Y4000 8) that shows that stretch (or equiva-lently, host galaxy class) is partially responsible. However at theshortest wavelengths (k ’ 3000Y3300 8, corresponding to ourUV1 diagnostic), additional effects are clearly important that arenot accounted for by the color-correction techniques in use incurrent cosmological programs. In addition to the newly foundintrinsic scatter at short wavelengths, new trends with phase arealso seen in the wavelengths of diagnostic features in this region.Redward of 40008, corrections for stretch or color work well

in normalizing SNe Ia; however, blueward of 40008, significantscatter remains even after such corrections are made. This maybe related to a change in the dominant source of opacity in SNe Ia.Redward of 3500Y40008 electron scattering opacity dominates,but at UVwavelengths a forest of overlapping lines is the dominantsource of opacity (seeFig. 1 of Hillebrandt&Niemeyer 2000). Elec-tron scattering is a continuous process involving well-understoodphysics, but line opacity depends sensitively on abundances, ion-ization states, and possibly non-LTE effects.Figure 16 illustrates this point by comparing the wavelength

dependence of the line and electron scattering opacity at maxi-mum light for a model by Kasen & Woosley (2007) that pro-vides a good match for a normal SN Ia (D. Kasen 2007, privatecommunication). The data in question refer to that at a depth of7000 km s!1, where intermediate-mass and Fe-peak material arewell mixed. A common feature in these models is the drop in lineopacity compared to the electron scattering opacity near 40008,as seen here. Thus, it is understandable that the emerging UV fluxis highly susceptible to changes in the line opacity (due to initialconditions and/or material synthesized during the explosion),while the optical and near-IR spectral behavior are dominated byelectron scattering opacity at this phase.What are the possible consequences of the above variations in

terms of the use of SNe Ia as probes of the expansion history? Inthe highest redshift surveys, including those proposed with futurefacilities, cross-color k-corrections are needed to estimate rest-frame light curves from the observations, typically undertaken in

Fig. 16.—Comparison of the wavelength-dependent line and electron scatteringopacity for a typical model in Kasen&Woosley (2007) at peak SN Ia brightness. Thedata refer to a depth of 7000 km s!1. Note the drop in line opacity with respect tothe electron scattering opacity near 4000 8. This behavior makes the emergentUV flux highly sensitive to changes in the line opacity, whereas the optical andnear-IR spectral regions are largely dominated by electron scattering opacity. [Seethe electronic edition of the Journal for a color version of this figure.]

ELLIS ET AL.66 Vol. 674Why the UV?

Ellis et al. 2008 (adopted from Kasen & Woosley 2007)

UV is dominated by line opacity

Optical is dominated by electron scattering

Metal content of SN should change the UV but leave the optical relatively unchanged

Using Spectra to Investigate EvolutionL112 OPTICAL SPECTRA OF TYPE Ia SUPERNOVAE Vol. 544

Fig. 1.—Heavily smoothed spectra of two high-z SNe Ia (SN 1999ff atand SN 1999fv at ) along with several low-z SN Ia spectraz p 0.455 z p 1.2

(SNe 1989B, 1992A, and 1981B) as well as SN IIb, Ib, and Ic spectra (SNe1993J, 1998dt [1983N in the UV], and 1994I). The date of the spectra relativeto B-band SN maximum is shown in parentheses after each object name.Specific features seen in SN 1999ff and labeled with a letter are discussed inthe text. The spectra have been scaled to the same flux level.

!23 mag. To eliminate contamination from the host galaxy, wescaled and subtracted the host galaxy spectrum from the SNspectrum. The high-z spectra have been smoothed with aSavitsky-Golay filter (Press et al. 1992), a polynomial fit thatpreserves line features better than boxcar smoothing. The spec-tra of SN 1999ff and SN 1999fv were heavily smoothed witha 70 and 180 A filter width, respectively, in the rest frame.For SN 1999ff we estimate the age based on its spectral

features (Riess et al. 1997) to be ! days. Photometric5! 2data of SN 1999ff will be presented by J. Tonry et al. (2000,in preparation) and will determine the true epoch of maximumbrightness. SN 1999fv was significantly fainter, and we wereunable to get a reliable spectral age for it, in part as a resultof the lack of available early-time spectra of low-z SNe Ia thatextend blueward of 4000 A.The redshift of the host galaxy of SN 1999ff is z p

, as determined by Balmer absorption lines in0.455! 0.001its spectrum, and we adopt this redshift for the SN. The hostgalaxy of SN 1999fv was not visible in our images, and tocalculate the redshift of this SN we cross-correlated the high-z spectrum with several low-z SN Ia spectra near maximumlight (as defined by the B-band maximum), the ages of whichare determined from their light curves. Correlating against asample of 31 spectra with ages between !8 and 5 days relativeto maximum, the redshift of SN 1999fv is forz p 1.17–1.22

the unsmoothed spectrum, somewhat lower than the initial es-timate of Tonry et al. (1999; ) made at the telescope.z p 1.23Using the same technique on the spectrum of SN 1999ff, cor-relating against a sample of 18 spectra with ages between !7and !3 days, we find , consistent with thez p 0.458! 0.006host galaxy redshift. We also include galaxy and M star spectraamong our templates for cross-correlation, which did not fitthe high-z spectra well.The redshift of SN 1999fv has the quoted uncertainty

(1.17–1.22) due to the low signal-to-noise ratio (S/N) of thespectrum and to its high redshift, placing key SN features nearbright sky lines at observed wavelengths !9000 A. In order tocheck the validity of our redshift estimates, we performed blindtests with fake spectra added to the two-dimensional frames.We redshifted each of 14 low-z SN Ia spectra by an arbitraryamount between and , scaled the flux levelz p 0.5 z p 1.3down to the signal of SN 1999fv in each 1000 s exposure,and added Poisson noise. We added the spectra to the two-dimensional frames (which are dominated by sky emission),convolving the signal with a Gaussian across several pixelsperpendicular to the dispersion, imitating seeing effects. Eachspectrum was then extracted, and we verified that its S/N agreedwith that of SN 1999fv. The spectrum was smoothed heavilywith a Savitsky-Golay filter to get an initial redshift estimateby eye and subsequently cross-correlated with a database oflow-z spectra to find a quantitative redshift estimate and un-certainty. We obtained good results for 12 of the 14 spectra,all of which matched the input redshift to within the error bars.For one spectrum we were unable to estimate the redshift, asthe signal was dominated by sky lines and SN features werenot clearly visible. For another spectrum we obtained twoequally likely redshifts, one of which was the correct inputredshift. In neither of these two cases were we led to believean incorrect redshift. The redshift range for SN 1999fv of

is at roughly 95% confidence level, based onz p 1.17–1.22the tests done. It is important to note that the difference inluminosity distance at and is 0.11 magz p 1.17 z p 1.22(5%), which is much smaller than our distance uncertainty.This redshift uncertainty therefore has negligible impact on ourability to constrain cosmology.The largest concern with using SNe Ia in cosmological tests

is a possible photometric difference between the low-z andhigh-z samples. Spectral comparisons of the distant and nearbySNe could show subtle evolutionary effects, if present. If thespectra do not show differences, this does not prove that thepeak luminosities of the SNe are identical, but it does buildconfidence that the two samples are similar. Spectra of normal,nearby SNe Ia at the same phase are quite homogeneous(Branch, Fisher, & Nugent 1993; Filippenko 1997; Riess et al.1997). As a qualitative comparison with our high-z spectra, wepresent in Figure 1 spectra of four low-z SNe Ia near the sameepoch: SN 1989B at both!7 and!1 days relative to maximum(Wells et al. 1994), SN 1992A at!1 day (Kirshner et al. 1993),and SN 1981B at maximum (Branch et al. 1983). The SN1989B spectra have been dereddened by E(B!V mag) p 0.32(Wells et al. 1994). The !1 day SN 1989B spectrum is acomposite of a Cerro Tololo Inter-American Observatory(CTIO) spectrum redward of 3300 A and an International Ul-traviolet Explorer (IUE) spectrum from !4 days blueward of3300 A. The SN 1992A spectrum is a composite of a CTIOspectrum redward of 3600 A and a Hubble Space Telescope(HST) spectrum blueward of 3600 A, taken 5 days after max-imum (Kirshner et al. 1993). For comparison, we also showSNe II, Ib, and Ic spectra: SN IIb 1993J (Filippenko, Matheson,

Coil et al. 2000

But most high-z SN spectra have much lower S/N than low-z SN spectra

Individual spectra look similar Foley et al. 2008b

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Coil et al. 2000



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Combine the Spectra!











Coil et al. 2000



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Combine the Spectra!Get Better Spectra!

Previous Studies Have Drawbacks

20 Ellis et al.

3000 3500 4000 4500 5000 5500Wavelength (Å)

0.0

0.5

1.0

1.5

2.0

F! +

Co

nst.

Mean high!z spectrum

Bootstrap sampled population

90% confidence

Local Nugent template

Mean high!z spectrum

Bootstrap sampled population

90% confidence

Local Nugent template

Maximum lightN (SNLS): 15

Fig. 5.— The mean high-redshift maximum light (e!ective day < ± 4 days) rest-frame UV SN Ia spectrum compared to the localaverage template of Nugent et al. (2002). Over-plotted in light grey are 100 bootstrap-resampled mean spectra drawn from the high-redshiftpopulation; the dotted lines show the region containing 90% of this distribution. The local template has been color-adjusted to match thehigh redshift data.

! ! ! ! ! !"#"

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Ellis et al. 2008

Foley et al. 2008aFoley et al. (ESSENCE) had large galaxy contamination

No individual spectra were very high S/N

Ellis et al. (SNLS) lacks a good low-z comparison

The wavelengths shown lacks the range to properly constrain the continuum

Motivation for Keck/SDSS Sample

0.10 0.15 0.20 0.25 0.30 0.35 0.40

Redshift

1

2

3

4

5

Nu

mb

er

of

SN

e I

a1. ESSENCE lacked individual high S/N spectra

2. SDSS had the photometry to remove galaxy light from SN spectra

3. This study requires a blue-sensitive spectrograph on a large-aperture telescope (LRIS)

4. z = 0.25 is the sweet spot:-SDSS can find them-Keck can get excellent spectra-Redshift rest-frame UV into optical

Foley et al. 2008b

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Motivation for Keck/SDSS Sample

0.10 0.15 0.20 0.25 0.30 0.35 0.40

Redshift

1

2

3

4

5

Nu

mb

er

of

SN

e I

a1. ESSENCE lacked individual high S/N spectra

2. SDSS had the photometry to remove galaxy light from SN spectra

3. This study requires a blue-sensitive spectrograph on a large-aperture telescope (LRIS)

4. z = 0.25 is the sweet spot:-SDSS can find them-Keck can get excellent spectra-Redshift rest-frame UV into optical

Foley et al. 2008b

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5. The mirror of the CTIO 4m (which ESSENCE used) was falling off, leaving us with a Keck night and no targets

Keck/SDSS Sample Demographics

0.10 0.15 0.20 0.25 0.30 0.35 0.40

Redshift

1

2

3

4

5

Nu

mb

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a

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More Luminous

Less Luminous

21 SNe Ia from z = 0.11 to 0.38.

Final sample has 15 SNe Ia

Mean z = 0.21

Mean phase = -0.6 d

Mean Delta = -0.13

Collaboration between SDSS-II SN and Foley/Filippenko

Targets selected to have z ~ 0.25 and phase ≤3 d

Keck spectra with SDSS light curves

Keck/SDSS Sample Demographics

0.10 0.15 0.20 0.25 0.30 0.35 0.40

Redshift

1

2

3

4

5

Nu

mb

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a

!! !" !# $ # " ! %&'()!*+,-'./0,('.&'1,)23'.)4.5,62-7-.89:

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More Luminous

Less Luminous

21 SNe Ia from z = 0.11 to 0.38.

Final sample has 15 SNe Ia

Mean z = 0.21

Mean phase = -0.6 d

Mean Delta = -0.13

Collaboration between SDSS-II SN and Foley/Filippenko

Targets selected to have z ~ 0.25 and phase ≤3 d

Keck spectra with SDSS light curves

– 36 –

!0.5 0.0 0.5 1.0 1.5

!

2

4

6

8

10

12

14

Num

ber

of

SN

e I

a

Fig. 20.— Histogram of the ! distribution of the SDSS SNe Ia from the first SDSS cosmo-logical analysis and the subsample presented in this paper (dashed histogram). The dashed

lines mark the regions of overluminous (! < !0.15), normal (!0.15 < ! < 0.3), andunderluminous (! > 0.3) objects, as defined by Jha et al. (2007) and shown in Figure 3.

Keck/SDSS Spectra

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12 of 15 Spectra in Final Sample

Higher S/N Spectra Show More Features

Foley et al. 2008c

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Foley et al. 2008b

Typical High-z SN Ia

Binned to 10 Å/pixel

Keck/SDSS SN Ia

Binned to 2 Å/pixel

SDSS Allows for Galaxy Correction1. Measure ugriz galaxy magnitudes from DR6

2. Use kcorrect to determine galaxy SED

3. Use synthetic magnitudes from SN spectrum to determine galaxy contamination level

4. Subtract galaxy SED from SN spectrum to get galaxy-subtracted SN spectrum

Host of SN 7147

SDSS Allows for Galaxy Correction1. Measure ugriz galaxy magnitudes from DR6

2. Use kcorrect to determine galaxy SED

3. Use synthetic magnitudes from SN spectrum to determine galaxy contamination level

4. Subtract galaxy SED from SN spectrum to get galaxy-subtracted SN spectrum !""" #""" $""" %""" &"""

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Foley et al. 2008c

Host of SN 7147

Low-z Composite For Comparison

Foley, Filippenko, & Jha 2008

Foley et al. 2008a

LickNugentSNLS

400 optical spectra from Lick/Keck over ~20 years

64 UV spectra from IUE/HST

! ! ! ! !"#"

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Fig. 1.— UV spectra of SNe Ia with t ≤ 1.6 d. The spectra are generally of low S/N,

but some features are apparent. In particular, there is an absorption feature at ∼3000A,attributed to Fe II (Branch & Venkatakrishna 1986), on either side of two peaks. Thesefeatures are seen in every spectrum with t < 3 weeks and of reasonable S/N.

Comparing Composite Spectra

Foley et al. 2008c

– 20 –

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Comparing Composite Spectra

Foley et al. 2008c

– 20 –

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UV Excess in Keck/SDSS

Objects

Comparing Individual Spectra

Foley et al. 2008c

3000 4000 5000 6000 7000Rest Wavelength (Å)

0.0

0.2

0.4

0.6

0.8

1.0

Rel

ativ

e f !

F275W

NUV

u g r

NUV

- g

F

275W

- g

u -

g

Comparing Individual Spectra

Foley et al. 2008c

3000 4000 5000 6000 7000Rest Wavelength (Å)

0.0

0.2

0.4

0.6

0.8

1.0

Rel

ativ

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– 25 –

! ! ! ! ! ! !

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Fig. 12.— Color-color diagrams for the individual Keck/SDSS SNe Ia, color coded in thesame way as in Figure 6. Specifically, the black, purple, green, and red lines correspond

to objects with no host galaxy identified in deep images, and no, low (< 8%), and high(> 8%) host-galaxy contamination, respectively. The color measurements are derived byconvolving the individual Keck/SDSS SN spectra with filter transmission curves. Since not

all of the spectra cover the entire wavelength range of each filter, there is a range of colorfor some points. The upper limits are derived by setting all flux for wavelengths shorter

than the shortest wavelength for a spectrum to zero. The lower limits are derived by settingall flux for wavelengths shorter than the shortest wavelength for a spectrum to the median

flux value of the bluest 100 A of the spectrum. The black cross with error bars in bothdirections indicates the colors measured for the low-redshift composite spectrum. The blackX indicates the SDSS composite spectrum.

NUV

- g

F

275W

- g

u -

g

– 45 –

0.6 0.8 1.0 1.2 1.4 1.6 1.8Observed Wavelength (µm)

0.0

0.2

0.4

0.6

0.8

1.0

Rel

ativ

e f !

z = 1.5

i F850LP F110W F160W

Fig. 24.— Composite SDSS spectrum shown in Figure 8 redshifted to z = 1.5. Overplotted

are (from blue to red) the transmission curves of the SDSS i, HST ACS F850LP , HST

NICMOS F110W , and F160W filters. The HST ACS F850LP filter is approximately the

same as the SDSS z band, the HST NICMOS F110W filter is a very wide filter including

the wavelengths covered by the z through J filters, and the HST NICMOS F160W filter is

slightly wider than the H-band filter.

Cosmological ImplicationsDon’t use the rest-frame U band, because if you do...

Assume that difference between spectra corresponds to luminosity difference

Measure the i-band difference for all redshifts

– 45 –

0.6 0.8 1.0 1.2 1.4 1.6 1.8Observed Wavelength (µm)

0.0

0.2

0.4

0.6

0.8

1.0

Rel

ativ

e f !

z = 1.5

i F850LP F110W F160W

Fig. 24.— Composite SDSS spectrum shown in Figure 8 redshifted to z = 1.5. Overplotted

are (from blue to red) the transmission curves of the SDSS i, HST ACS F850LP , HST

NICMOS F110W , and F160W filters. The HST ACS F850LP filter is approximately the

same as the SDSS z band, the HST NICMOS F110W filter is a very wide filter including

the wavelengths covered by the z through J filters, and the HST NICMOS F160W filter is

slightly wider than the H-band filter.

Cosmological ImplicationsDon’t use the rest-frame U band, because if you do...

Assume that difference between spectra corresponds to luminosity difference

Measure the i-band difference for all redshifts

– 46 –

0.0 0.5 1.0 1.5 2.0Redshift

!0.2

0.0

0.2

0.4

0.6

0.8

i!ban

d D

iffe

rence

(m

ag)

Fig. 25.— Di!erence between observer-frame i-band magnitudes for the artificially redshiftedSDSS and low-redshift composite spectra. The measurement for each spectrum is normalized

to the magnitude at z = 0.13, the lowest redshift with complete coverage of the i band for

the SDSS composite spectrum. The grey region is the uncertainty, which is a combination

of the boot-strap sampling errors and a 0.03 mag error in the photometry. Also plotted are

the di!erences from the measured distance moduli of the two z > 1.5 SNe Ia from Riesset al. (2007) compared to the expected distance moduli for the redshifts of those objects in

the "CDM concordance cosmology. At z = 1.5 the di!erence in the i-band measurements is

!0.5 mag.

Hubble residuals for highest redshift SNe Ia match this difference

Hubble residuals forz > 1.5 SNe Ia

Potential Biases

High-z sample is not representative

Preferentially selecting blue objects

Spectral reductions, galaxy subtraction, making the composite spectrum are wrong

We chose the wrong RV

Low-z UV spectra are not representative

Malmquist bias?

Incorrect treatment of dust/intrinsic colors

Potential Biases that don’t cause this

Potential Biases that could cause this

Conclusions and Future Work

There is a UV difference for our two SN samples

It is consistent with SN evolution

It may be the result of a poor low-z sample or possibly a Malmquist bias

Do not use the U (and B?)band for cosmology

We need more high-z spectra to determine if there is a Malmquist bias

We desperately need more low-z UV spectra

We can check these results with photometry (more tomorrow)

The Keck/SDSS SN Samplesdss2008.uchicago.edu/depot/foley-ryan.pdf · Fig. 8.Ñ (top panel):...

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