A chemical signature from fast-rotating low-metallicity massive stars: ROA 276 in omega Centauri

TL;DR

This study analyzes the chemical composition of star ROA 276 in omega Centauri, providing evidence that fast-rotating low-metallicity massive stars contributed to early Universe heavy element production.

Contribution

It presents the first detailed chemical signature linking a specific star to nucleosynthesis in fast-rotating massive stars at low metallicity.

Findings

ROA 276 has an unusually high [Sr/Ba] ratio.

The star's abundance pattern matches models of 20 Msun fast-rotating stars.

No other nucleosynthetic source explains the neutron-capture elements.

Abstract

We present a chemical abundance analysis of a metal-poor star, ROA 276, in the stellar system omega Centauri. We confirm that this star has an unusually high [Sr/Ba] abundance ratio. Additionally, ROA 276 exhibits remarkably high abundance ratios, [X/Fe], for all elements from Cu to Mo along with normal abundance ratios for the elements from Ba to Pb. The chemical abundance pattern of ROA 276, relative to a primordial omega Cen star ROA 46, is best fit by a fast-rotating low-metallicity massive stellar model of 20 Msun, [Fe/H] = -1.8, and an initial rotation 0.4 times the critical value; no other nucleosynthetic source can match the neutron-capture element distribution. ROA 276 arguably offers the most definitive proof to date that fast-rotating massive stars contributed to the production of heavy elements in the early Universe.

Tables

Table 1. Table 1: Stellar parameters for the program stars.

Star	$T_{\mathrm{eff}}$	$\log g$	${\xi }_t$	[Fe/H]	$T_{\mathrm{eff}}$	$\log g$
	(K)	(cgs)	(km s-1)	(dex)	(K)	(cgs)
	Spectroscopic				Photometric
ROA 276	4125	0.70	1.75	$-$1.30	4130	0.79
ROA 46	4075	0.20	2.40	$-$1.72	4024	0.37

Table 2. Table 2: Line list, equivalent widths, and abundances.

Species	Wavelength	LEP	$\log gf$	EW (ROA 276)	EW (ROA 46)	$\log ϵ$(X) (ROA 276)	$\log ϵ$(X) (ROA 46)	Source
	(Å)	(eV)		(mÅ)	(mÅ)	(dex)	(dex)
CH	4270 - 4330			syn	syn	6.68	6.08	1
O I	6300.31	0.00	$-$9.75	69.6	69.7	8.00	7.60	2
O I	6363.78	0.02	$-$10.25	36.9	…	8.09	…	3
Na I	5682.65	2.10	$-$0.67	64.8	…	4.82	…	2
Na I	5688.22	2.10	$-$0.37	87.4	68.7	4.86	4.55	2

Table 3. Table 3: Chemical abundances of ROA 276 and the comparison star ROA 46.

\@alignment@align    Species        \@alignment@align

$\log ϵ$

${\sigma }_{\log ϵ}$

[Fe/H]

${\sigma }_{\left[\mathrm{Fe}/H\right]}$

$\log ϵ$

${\sigma }_{\log ϵ}$

[Fe/H]

${\sigma }_{\left[\mathrm{Fe}/H\right]}$

$\Delta$[Fe/H]

ROA 276

ROA 46

ROA 276 $-$

ROA 46

Fe I

6\@alignment@align.20

0.01

$-$1\@alignment@align.30

0.08

5\@alignment@align.78

0.02

$-$1\@alignment@align.72

0.08

0\@alignment@align.42

Fe II

6\@alignment@align.21

0.04

$-$1\@alignment@align.29

0.15

5\@alignment@align.80

0.03

$-$1\@alignment@align.70

0.13

0\@alignment@align.41

$\log ϵ$

${\sigma }_{\log ϵ}$

[X/Fe]

${\sigma }_{\left[X/\mathrm{Fe}\right]}$

$\log ϵ$

${\sigma }_{\log ϵ}$

[X/Fe]

${\sigma }_{\left[X/\mathrm{Fe}\right]}$

$\Delta$[X/Fe]

ROA 276

ROA 46

ROA 276 $-$

ROA 46

C (CH)

6\@alignment@align.68

0.20

$-$0\@alignment@align.45

0.24

6\@alignment@align.08

0.20

$-$0\@alignment@align.63

0.24

0\@alignment@align.18

O I

8\@alignment@align.04

0.05

0\@alignment@align.65

0.18

7\@alignment@align.60

0.20

0\@alignment@align.63

0.22

0\@alignment@align.02

Na I

4\@alignment@align.76

0.08

$-$0\@alignment@align.19

0.13

4\@alignment@align.55

0.20

0\@alignment@align.03

0.21

$-$0\@alignment@align.22

Mg I

6\@alignment@align.85

0.02

0\@alignment@align.54

0.13

6\@alignment@align.34

0.20

0\@alignment@align.46

0.22

0\@alignment@align.08

Ca I

5\@alignment@align.54

0.03

0\@alignment@align.50

0.14

5\@alignment@align.00

0.04

0\@alignment@align.38

0.13

0\@alignment@align.12

Sc II

1\@alignment@align.48

0.06

$-$0\@alignment@align.37

0.14

1\@alignment@align.49

0.06

0\@alignment@align.06

0.13

$-$0\@alignment@align.43

Ti I

4\@alignment@align.12

0.02

0\@alignment@align.47

0.13

3\@alignment@align.70

0.03

0\@alignment@align.47

0.13

0\@alignment@align.00

Ti II

4\@alignment@align.21

0.05

0\@alignment@align.56

0.14

3\@alignment@align.69

0.03

0\@alignment@align.46

0.12

0\@alignment@align.10

Cr I

4\@alignment@align.45

0.06

0\@alignment@align.11

0.11

4\@alignment@align.03

0.12

0\@alignment@align.10

0.16

0\@alignment@align.01

Cr II

4\@alignment@align.35

0.12

0\@alignment@align.01

0.18

4\@alignment@align.12

0.07

0\@alignment@align.20

0.17

$-$0\@alignment@align.19

Mn I

3\@alignment@align.71

0.03

$-$0\@alignment@align.42

0.10

3\@alignment@align.36

0.06

$-$0\@alignment@align.35

0.10

$-$0\@alignment@align.07

Co I

3\@alignment@align.91

0.04

0\@alignment@align.22

0.17

3\@alignment@align.25

0.20

$-$0\@alignment@align.02

0.21

0\@alignment@align.24

Ni I

5\@alignment@align.27

0.03

0\@alignment@align.35

0.08

4\@alignment@align.56

0.03

0\@alignment@align.06

0.09

0\@alignment@align.29

Cu I

3\@alignment@align.56

0.02

0\@alignment@align.67

0.17

1\@alignment@align.96

0.06

$-$0\@alignment@align.51

0.16

1\@alignment@align.18

Zn I

4\@alignment@align.73

0.18

1\@alignment@align.47

0.20

2\@alignment@align.98

0.09

0\@alignment@align.14

0.16

1\@alignment@align.33

Rb I

3\@alignment@align.15

0.03

1\@alignment@align.93

0.17

1\@alignment@align.22

0.20

0\@alignment@align.42

0.23

1\@alignment@align.51

Sr I

2\@alignment@align.90

0.03

1\@alignment@align.32

0.17

0\@alignment@align.78

0.20

$-$0\@alignment@align.37

0.21

1\@alignment@align.69

Y II

2\@alignment@align.21

0.08

1\@alignment@align.30

0.14

0\@alignment@align.33

0.08

$-$0\@alignment@align.17

0.15

1\@alignment@align.47

Zr I

3\@alignment@align.12

0.13

1\@alignment@align.84

0.16

1\@alignment@align.16

0.08

0\@alignment@align.30

0.10

1\@alignment@align.54

Zr II

3\@alignment@align.02

0.41

1\@alignment@align.74

0.43

1\@alignment@align.29

0.06

0\@alignment@align.43

0.17

1\@alignment@align.31

Mo I

1\@alignment@align.92

0.20

1\@alignment@align.34

0.22

0\@alignment@align.11

0.20

$-$0\@alignment@align.05

0.22

1\@alignment@align.39

Ba II

1\@alignment@align.21

0.09

0\@alignment@align.33

0.14

0\@alignment@align.56

0.12

0\@alignment@align.09

0.14

0\@alignment@align.24

La II

$-$0\@alignment@align.12

0.09

0\@alignment@align.08

0.13

$-$0\@alignment@align.66

0.07

$-$0\@alignment@align.04

0.10

0\@alignment@align.12

Ce II

0\@alignment@align.07

0.09

$-$0\@alignment@align.21

0.13

$-$0\@alignment@align.32

0.07

$-$0\@alignment@align.18

0.13

$-$0\@alignment@align.03

Pr II

$-$0\@alignment@align.62

0.16

$-$0\@alignment@align.04

0.19

$-$1\@alignment@align.25

0.00

$-$0\@alignment@align.25

0.15

0\@alignment@align.21

Nd II

0\@alignment@align.22

0.05

0\@alignment@align.10

0.10

$-$0\@alignment@align.26

0.04

0\@alignment@align.04

0.07

0\@alignment@align.06

Sm II

$-$0\@alignment@align.64

0.06

$-$0\@alignment@align.30

0.13

$-$0\@alignment@align.78

0.05

$-$0\@alignment@align.02

0.11

$-$0\@alignment@align.28

Eu II

$-$0\@alignment@align.76

0.20

0\@alignment@align.02

0.22

$-$1\@alignment@align.17

0.20

0\@alignment@align.03

0.20

$-$0\@alignment@align.01

Pb I

0\@alignment@align.70

0.20

0\@alignment@align.25

0.23

0\@alignment@align.45

0.20

0\@alignment@align.42

0.24

$-$0\@alignment@align.17

Full text

stars: abundances — stars: Population II — globular clusters: individual: $\omega$ Centauri

††facilities: Magellan:Clay (MIKE)

1 Introduction

Numerical simulations predict that low-metallicity stars that formed in the early Universe were massive, compact, and rotated near their critical velocities where gravity is balanced by centrifugal forces (Bromm & Larson, 2004; Stacy et al., 2011). Nucleosynthesis in these fast-rotating low-metallicity massive stars (hereafter spinstars) differs considerably from their non-rapidly-rotating counterparts (Meynet et al., 2006; Hirschi, 2007; Pignatari et al., 2008; Frischknecht et al., 2012, 2016; Maeder & Meynet, 2012). Since these massive stars have long since died, confirmation of their existence can be obtained by identifying their unique chemical signature in the abundance patterns of subsequent generations of Milky Way stars (Frebel & Norris, 2015; Maeder et al., 2015).

One chemical signature of spinstars comes from nitrogen abundances in metal-poor halo stars which require primary production (Spite et al., 2005). While spinstars can naturally achieve such nucleosynthesis, hydrogen ingestion in massive stars (Pignatari et al., 2015) and intermediate-mass and super asymptotic giant branch (AGB) stars (Karakas, 2010; Doherty et al., 2014) may also be responsible for nitrogen production in the early Universe.

Another possible observational signature of spinstars comes from neutron-capture elements. The scatter in Sr and Ba abundances in low-metallicity halo stars can be explained by spinstars (Cescutti et al., 2013), but measurements of other neutron-capture elements (e.g., Y, Zr, La), when available, are also compatible with massive AGB stars (Fishlock et al., 2014). Chiappini et al. (2011) reported unusually high abundances for the elements Sr, Y, Ba, and La in the bulge globular cluster NGC 6522, consistent with yields from spinstars. Those measurements, however, have since been revised downwards and could also be explained by AGB stars (Barbuy et al., 2014; Ness et al., 2014). The unmistakable signature among the neutron-capture elements from spinstars has yet to be seen within an individual star.

2 Target Selection and Observations

$\omega$ Centauri is the most massive star cluster in our Galaxy. In contrast to the majority of Milky Way globular clusters, $\omega$ Cen exhibits a number of peculiar features including a broad range in abundances for iron and slow neutron-capture process, or $s$-process, elements (Norris & Da Costa, 1995). The distribution and evolution of the $s$-process element abundances in $\omega$ Cen are consistent with a dominant contribution from 1.5 - 3 M*⊙* AGB stars (Smith et al., 2000).

There are two stars in $\omega$ Cen, however, that exhibit peculiar abundance ratios of Sr and Ba (Stanford et al., 2006, 2010); the red giant ROA 276 with $V$ = 12.37 and the main sequence star 2015448 with $V$ = 18.22. Both objects have high Sr and low Ba abundances, consistent with predictions of neutron-capture nucleosynthesis in spinstars (Frischknecht et al., 2012, 2016).

To further examine these unusual abundance patterns, we obtained a high-resolution optical spectrum for the red giant ROA 276 and a comparison star ROA 46 ($V$ = 11.54) using the Magellan Inamori Kyocera Echelle spectrograph (Bernstein et al., 2003) at the 6.5m Magellan Clay Telescope on 2007 June 22-23. Both stars have proper motions and radial velocities consistent with cluster membership (Bellini et al., 2009). The total exposure time was 10 min per target. We used the 0.5″ slit to achieve a spectral resolution of $R$ = 56,000 and $R$ = 44,000 in the blue and red arms, respectively. One dimensional, wavelength calibrated, continuum normalized spectra were produced from the raw spectra using iraf222IRAF is distributed by the National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy (AURA) under a cooperative agreement with the National Science Foundation. and the mtools333www.lco.cl/telescopes-information/magellan/instruments/mike/iraf-tools/iraf-mtools-package package. The signal-to-noise ratio (S/N) for both stars was roughly 80 per pixel near 6000 Å and 40 per pixel near 4500 Å. The spectra have approximately 3.5 pixels per resolution element.

3 Stellar Parameters and Chemical Abundances

The stellar parameters were determined from a traditional spectroscopic approach following the procedure outlined in Yong et al. (2014). Equivalent widths (EWs) were measured using routines in iraf and daospec (Stetson & Pancino, 2008), and there was good agreement between the two approaches. Weak (EW $<$ 10 mÅ) and strong (EW $>$ 150 mÅ) lines were removed from the analysis. Abundances were derived using the EW, one-dimensional local thermodynamic equilibrium (LTE) model atmospheres with [$\alpha$/Fe] = +0.4 (Castelli & Kurucz, 2003), and the LTE stellar line analysis program moog (Sneden, 1973). The version of moog that we used includes a proper treatment of Rayleigh scattering (Sobeck et al., 2011). The effective temperature ($T_{\rm eff}$), surface gravity ($\log g$), and microturbulent velocity ($\xi_{t}$), were obtained by enforcing excitation and ionization balance for Fe lines (see Table 1). The uncertainties in $T_{\rm eff}$, $\log g$, and $\xi_{t}$ are 50 K, 0.2 dex, and 0.2 km s*-1*, respectively. The standard deviation for Fe i lines was 0.19 dex (ROA 276) and 0.16 dex (ROA 46), and we adopted an uncertainty in the model atmosphere of [m/H] = 0.2 dex.

Stellar parameters can also be derived from a photometric approach. $T_{\rm eff}$ can be estimated from color-temperature relations based upon the infrared flux method (Blackwell & Shallis, 1977; Ramírez & Meléndez, 2005). We used $BVRIJHK$ photometry (Bellini et al., 2009; Skrutskie et al., 2006) and adopted a reddening of $E(B-V)$ = 0.12 (Harris, 1996, 2010 edition). The surface gravity can be determined assuming the photometric $T_{\rm eff}$, a distance modulus $(m-M)_{V}$ = 13.94 (Harris, 1996, 2010 edition), bolometric corrections from Alonso et al. (1999), and a mass of 0.8 M*⊙*. $T_{\rm eff}$ and $\log g$ obtained from the spectroscopic and photometric approaches are in good agreement when considering the estimated uncertainties (see Table 1).

Elemental abundances were derived using moog for individual lines based on the EW or from spectrum synthesis following Yong et al. (2014). Examples of synthetic spectra fits for representative lines of selected elements are given in Figure 1. Aside from the 4057.81 Å Pb i line, given the S/N of the blue spectra we analyzed lines redward of 4317.31 Å. We present our line list, EWs, and abundance measurements in Table 3. Solar abundances were taken from Asplund et al. (2009) and the sources of the $gf$ values can be found in Table 3.

Uncertainties in chemical abundances were obtained by repeating the analysis and varying the stellar parameters, one at a time, by their uncertainties. These four error terms were added, in quadrature, to obtain the systematic uncertainty. We replaced the random error (s.e.logϵ) by max(s.e.logϵ, 0.20/$\sqrt{N_{\rm lines}}$) where the second term is what would be expected for a set of $N_{\rm lines}$ with a dispersion of 0.20 dex. The total error was obtained by adding the random and systematic errors in quadrature. Chemical abundances and their errors are presented in Table 3.