The Remarkable Chemical Compositions of Blue Metal-Poor Stars George Preston on behalf of Chris Sneden friends & collaborators George Preston (Carnegie Observatories) John Cowan (University of Oklahoma) Ian Thompson, Steve Shectman (Carnegie Observatories

Outline of the talk What are blue metal-poor (= BMP) stars? Use of binary fractions to resolve BMPs into blue straggler and intermediate-age populations Fundamental differences between blue stragglers in globular clusters and the halo field Chemical compositions: general results A new spectroscopic study: binaries versus single stars Present & future observational/theoretical opportunities

Identifying field BMP stars: Galactic disk star color-color relation BMP domain can be populated by metal-poor MS stars. It is almost empty in the solar neighborhood Preston et al Data points are the B8-F0 stars in the Bright Star Catalog [Fe/H] ~ 0 [Fe/H] ~  1 de-blanketing (blue) vectors [Fe/H] ~  3  

BMP domain is a region in which isochrones for a wide range of ages and metallicities overlap in a tangled mess. Preston & Sneden 2000 Isochrones of various [Fe/H] values and ages overlap in the BMP star domain Isochrones in the U-B versus B-V plane are from Green et al. (1987) Revised Yale Isochrones MS isochrones SG isochrones Turnoffs for: [Fe/H] =  2.2 ages 3 7,10 Gy

BMP stars were readily identified by UBV photometry of stars found in the HK objective-prism survey Preston et al BMP stars “HK” Survey: Beers et al. 1985, 1992 MS [Fe/H]=0 MS [Fe/H]=  1 RHB Metal poor stars near turnoff BHB

Blue metal-poor stars of the halo field have many physical characteristics of the blue stragglers first identified in globular clusters, but there is a problem: THERE ARE FAR TOO MANY OF THEM IN THE HALO FIELD! NGC 288 (Kaluzny 1996) Common distance makes identification easy MS RGB BHB RHB AGB MS turnoff BMP

Resolution of the problem lay in a radial velocity survey: a surprisingly high proportion of BMP stars are members of spectroscopic binaries Preston & Sneden km/s  < 1.5 km/s: 1/18 stars with orbits  > 1.5 km/s: 41/43 stars with orbits 1.5 km/s is an observational limit; more low-amplitude binaries may be hidden in the RV errors. >60% of BMP’s are binaries with atypically long orbital periods and small mass functions. Standard deviation of a RV measurement (km/s)

Typical orbital solutions based on radial velocity variations in the BMP sample Preston & Sneden 2000 >60% of BMP’s are binaries with unusually long orbital periods and unusually small mass functions. f(m)  (K 1 ) 3 P

We use binary fractions to resolve blue metal-poor stars into blue straggler and intermediate-age components. xxxxxxxxxxxxxxx n

To estimate the BS fraction of BMP  n(BS)/n(BMP) f BMP, f BS, and f IA must be “good numbers”. COMMENTS f BMP  Reliability limited only by accuracy of RV’s & duration of survey f BS  Adopt Duquennoy & Mayor n(P) for primordial blue stragglers.  Assume that 13% of primordial blue straggler binaries with P < 5 d have merged.  Remainder (87%) must be mass-transfer binaries. f IA  Adopt 0.15 as “universal binary fraction” for P<4000 d from: f Disk = 0.15 Duquennoy & Mayor 1991 f Halo = 0.14 Latham et al 1998

Only this 13% of binaries with P  4000d can merge in a Hubble time (Vilhu, O. 1982, A&Ap, 109, 17) This is why we adopt f BS  0.87 radial velocity binaries visual binaries c.p.m. binaries {

To estimate the BS fraction of BMP  n(BS)/n(BMP) f BMP, f BS, and f IA must be “good numbers”. COMMENTS f BMP  Reliability limited only by accuracy of RV’s & duration of survey f BS  Adopt Duquennoy & Mayor n(P) for primordial field blue stragglers.  Assume that 13% of primordial blue straggler binaries with P < 5 d have merged.  Remainder (87%) must be mass-transfer binaries. f IA  Adopt 0.15 as “universal binary fraction” for P<4000 d from f Disk = 0.15 Duquennoy & Mayor 1991 f Halo = 0.14 Latham et al 1998

the blue straggler fraction of BMP stars is f BMP f IA f BS Binary Fraction (f) n BS /n BMP = (f BMP  f IA )/(f BS  f IA ) = 0.62 More than half of the blue metal-poor stars are blue stragglers. BUT Inserting our adopted binary fractions for the three populations

Halo field blue stragglers (FBS) are a different breed. S = Specific frequency = (BMP)/ (HB) BMP = 300 kpc -3 S BMP = 6.7 HB = 45 kpc -3  BS ~  10 8 y  M(parent pop) S HFBS = 0.62*6.7 = 4.2 Specific frequency of halo field blue stragglers exceeds specific frequency of blue stragglers in globular clusters by factor 10:

Specific frequency of blue stragglers in globular clusters (~0.4) is one order-of-magnitude smaller than value in halo field (~4). increasing cluster mass Blue stragglers occur more frequently in less massive, loosely-bound clusters                halo field blue stragglers

Mateo, Harris, Nemec, & Olszewski 1990, Astronomical Journal, 100, 469 Mateo et al (1990) used estimates of merger time-scale (~5E+8 y) and blue straggler lifetime (  7E+9 y) to conclude that the specific frequency of blue stragglers in NGC 5466 can be explained entirely by mergers of the cluster population of W Uma systems.

Mapelli et al. (2004) simulations of 47 Tuc data confirm that mergers produce the observed S GCBS outside of the core. Observed distribution (Ferraro et al 2004) Collisional formation only Collisions plus binary mergers

TO SUMMARIZE We use binary fractions to resolve BMPs into two populations: 40% intermediate-age, metal-poor stars (IA) 60% old metal-poor blue stragglers (FBS) A small fraction of HFBS are formed by merger of close pairs. The rest must be formed by McCrea mass transfer, because there are no collisions in the halo. GCBS are formed primarily by collisions (in core) and mergers (everywhere) of the small portion (10%) of primordial binaries that survive disruption by encounters. By this reasoning we understand why the specific frequency of HFBS exceeds that of GCBS by an order-of-magnitude.

Abundance analysis of Las Campanas high resolution spectra (R ~ 25,000) of BMP stars We analyzed summed spectra. Individual frames used to search for velocity variations have too small S/N to be employed in abundance work. Preston & Sneden 2000 Vsini=40 km/s [Fe/H]=-2.30

The abundance analysis Use Fe-peak lines to derive atmosphere parameters Interpolated model atmospheres from Kurucz’s ATLAS grid Standard LTE analysis with the MOOG code T eff from Fe I abundances with excitation potential log g from Fe I versus Fe II abundances v t from Fe I abundances with EW Basic results: (1) overall metallicities, (2) Teff’s in K range, (3) main-sequence gravities Limited # of lines → abundances of 8 elements

Abundances from the Las Campanas BMP high resolution survey Normal results (  ) compared to other Pop II halo stars: (1) Mg, Ca, Ti ↑ (2) Mn ↓ (3) Sr, Ba: large  at lowest [Fe/H] Preston & Sneden 2000 Open circles denote stars with v e sini > 24 km/s  lower accuracy

Orbital parameters for ordinary binary stars... and for Carbon-rich & s process-rich binaries Preston & Sneden 2000 The fraction of binaries with P > 25d & e < 0.15 is small. The fraction of binaries with P > 25d & e < 0.15 is larger. Mass transfer during post-MS evolution circularizes orbits  25 d 

Orbital parameters for ordinary binary stars … and for BMP stars Preston & Sneden 2000 The fraction of binaries with P > 25d & e < 0.15 is small. The fraction of binaries with P > 25d & e < 0.15 is larger. Mass transfer during post-MS evolution circularizes orbits  25 d 

Followup high resolution BMP study 5 BMP binaries, 5 BMP RV-constant stars Las Campanas echelle, with new CCD detector → higher S/N data Original goal: a comparative Li abundance study Oops!: Li undetected in all 10 stars! Much more interesting: look at the 3/5 stars in each group with [Fe/H] < -2

Radial velocities of the low metallicity BMP star sample: Some of them are binaries and others are not. RV-constant starsBinary stars

Individual and mean spectra of low metallicity RV-constant and binary stars High excitation O I lines are somewhat stronger in the binaries.

 -capture elements are ~ normal in the whole BMP sample Original BMP sample: Preston & Sneden 2000 neutron exposure ~ constant Upper-envelope for n-capture elements declines in the binaries as if neutron exposure is ~ constant for all [Fe/H].

Carbon species in the spectra of BMP RV- constant stars and binary CS CH & C I respond VERY differently to changes in temperature & gravity! The “non-variable” spectrum is the mean of three stars

Neutron-capture species in BMP RV- constant stars and binary CS (another lead-rich star) Note similarity of lines for Fe-peak elements Preston & Sneden 2000 The “non-variable” spectrum is the mean of three stars

Mean abundances in the low-metallicity binary and RV-constant groups Large  values for C, Sr, & Ba in the binaries indicate real star-to-star differences

Abundances in CS and the RV-constant stars Abundances of C, O, and n-capture elements are new; other abundances are from Preston & Sneden 2000

We know of several very lead-rich stars Abundances are normalized to CS at Ba, or La, or both Normalizations are simple vertical shifts Sneden et al. 2003

s-process predictions versus abundances in lead-rich stars Mean observed abundances are computed after normalizations Neutron/seed ratio is the main variable in the theoretical computations Arbitrary normalization between theory and observation at Ba & La

Pb-rich stars: a unique abundance signature? Domain of the enhanced r-process metal-poor stars Domain of the large s-process, lead-rich stars Halo sample w [Fe/H]<-1.5

Evolutionary states of known lead-rich stars Here is a new one: CS , an accidental discovery in a survey of 25 metal-poor red horizontal-branch (RHB) stars Preston et al. 2004

Most detailed n-capture abundance pattern of any lead-rich star? Preston et al Other stuff about CS : (1) It is relatively carbon-rich (like all other lead stars) (2) It is an RR Lyrae star (P=0.59 d) for heaven’s sake! (3) Is it in a binary? It ought to be! Time will tell.

Recapitulate BMP stars are in the “wrong” HR diagram place for metal-poor main sequence stars 2/3 of BMP stars are BS binaries. 5/7 very metal-poor BMP binaries are rich in s-process products → AGB mass transfer Companion stars must now be compact objects Pb discovered in one star; others must exist The Pb-rich turn-off stars must have experienced AGB mass transfer? How does this happen? Question: How can mass transfer be so efficient in (now) widely separated pairs?

The End