extreme helium stars: a panorama

Simon Jeffery - Armagh Observatory

extreme helium stars: a panorama

Armagh ObservatorySimon Jeffery, A Ahmad, T Sahin,

N Behara*, V Woolf*, K Gourgouliatos*, R Starling*, P Montanes Rodriguez*

* now elsewhere

H Saio (Sendai), D Pollacco (QUB), P W Hill (St Andrews)

D Lambert (UT), G Pandey, N K Rao (IIA, Bangalore)

The British Council, JSPS, PPARC

outline

• Basic properties• Abundances• Model atmospheres• Pulsation• Emission • Evolution• Links to other classes

Basic properties

• 1998 … more stars now measured …

Photospheric Abundances

Since 1995, analyses of optical and/or UV spectra to give abundances for 17 extreme helium stars, Pandey et al. (2006)

a) Proxies for metallicity (Ni,Mn,Cr,Fe)b) Overabundant light elements (Mg,Si,S,..?)

Photospheric Abundances II

c) C always overabundantd) N overabundant (=C+N+O)e) O ranges from overabundant to normal?ef) s-process “ “ “ “

Model atmospheres

• 1995 (as used by Jeffery et al., Pandey et al.)

– LTE – ATLAS6 continuous opacities– + Peach tables for C and N– ATLAS6 odfs

• 2007– OP continuous opacities (Behara & Jeffery 2006)– Opacity sampling (Behara & Jeffery 2006,

Przybilla et al. 2005)

– nLTE line formation (Przybilla 2005)

New opacities

Behara & Jeffery 2006

H-poor

H-rich

Consequences

• Behara & Jeffery applying models to IUE/HST• Significant change on Teff scale (~ 2000 K)• Impact on abundances not yet assessed

• Przybilla et al. (2006) used LTE atmospheres with OS and nLTE line formation and found

• Better agreement between Tflux and Tion • Better He line cores• Better W agreement in Balmer lines (V652 Her)• Substantial improvement in consistency of metal lines

• Low optical depth in H-def atmospheres currently remains a challenge for routine analysis of EHe stars.

• However, basic result for abundances unlikely to change by a huge amount … (is this tempting fate? … see Gustafsson’s review!)

pulsation

pulsations

Low Luminosity, Z bumpHigh Luminosity, High TeffHigh Luminosity, Low Teff

Low Luminosity, Z bump

• V652 Her– New hires data and model

– Baade mass: 0.58 M (Jeffery et al. 2001)

– N-rich

• LSS 3184 - discovered 1994 (Kilkenny et al.)– Baade mass: 0.4 M (Kilkenny et al. 1999, Woolf

& Jeffery 2002)– C-rich (Drilling et al. 1998)

V652 Her

WHT+ISIS 1998

Jeffery, Woolf & Pollacco (2001)

SAAO 1.0m StAP+CCD V-band

data from Kilkenny (priv. comm.)

P = 0.108 days

V652 Her the pulsating helium star

Pulsating Extreme helium star (Berger & Greenstein 1953, Landolt 1973)P=0.108 days

Radius measurement from Baade’s method (Jeffery et al. 2001)<R>=2.31±0.02R M = 0.59±0.18M

Period change (Kilkenny & Lynas-Gray 1982 - 1996) R’/R ~ 2 .10-4 yr-1 , R’’’)

Pulsation models (Saio 1983 - 1995, Fadeyev & Lynas-Gray 1996, Montanes Rodriguez & Jeffery 2001) 1993: OPAL and OP opacities => Z-bump opacity driving

H-poor, N-rich (Jeffery et al. 1999)1% H, [Fe/H]=0, N-rich, C and O poorlog g = 3.7 +/– 0.1Teff = 24 500 +/– 500 K

pulsations

Low Luminosity, Z bumpHigh Luminosity, High TeffHigh Luminosity, Low Teff

High Luminosity, Radial, Strange Mode

• HD168476 = PV Tel– 7 d ?

• BD-1 3438 = NO Ser– 5 - 7 d

• LS IV-1 2– 11 d

• BD+1 4381 = FQ Aqr– 19 - 24 d ?

• P Teff-1

SAAO 0.5m: FQ Aqr: P ~ 20d ?

Kilkenny et al. (1999) “…the apparent presence of several periods but, if real, none seems to persist for more than one season.”

pulsations

Low Luminosity, Z bumpHigh Luminosity, High TeffHigh Luminosity, Low Teff

High Luminosity, Non-Radial, Strange Mode

• BD-9 4395 = V2209 Oph– 4 - 9 d

• HD 160641 = V2076 Oph– 0.7 - 1.1 d

• LSS 5121 – ~ few hours

V2076 Oph

Wright et al. (2006), “ … conventional Fourier analysis … fails to reveal coherent frequencies … could be a result of random variations …”

Lynas-Gray et al. (1987).

Monitoring Programs

• SAAO + 0.5m (Kilkenny et al. 1999)– FQ Aqr : BD+1 4381– NO Ser : BD-1 3438

• SuperWASP (Sneak Preview)– V652 Her: BD+13 3224– FQ Aqr : BD+1 4381 – V821 Cen? : HD124448

SuperWASP: V652 Her: P = 0.108d

SAAO 1.0m

SuperWASP: FQ Aqr: P ~ 20d ? mag

mag

mag

day

day

day-1

day-1

Are FQ Aqr pulsations periodic?

SuperWASP: HD124448: non-variable

V2076 Oph: HIRES

Spectroscopy

• 2005 August• AAT/UCLES• 4 nights

V2076 Oph: HIRES

Spectroscopy

• Yellow: mean• Black:

residual• Red: model

V2076 Oph: rms

• Black: mean spectrum• Red: rms residual spectrum x10

– HeI– HeII – 4650 complex– CII 4267

V2076 Oph: line profile variations

Yellow: meanBlack: residualRed: model

CII 4267

OII, CIII

emission

emission lines in V2076 Oph

1986 Heber:

4486+4504 Also seen in

Of stars, unidentified

2001 Werner & Rauch:

Identified as highly excited transition of S IV

emission lines in V2076 Oph

2005 AAT:

C II, N II, Si III and a few unidentified em lines

Other lines seen in Of stars are not seen, HeII, NIII

Red spectra ?

evolution

CO + He WD merger

He+He white dwarf formed from binary star evolution (observed)

orbit decays through gravitational, tidal and magnetic interaction

less massive WD disrupted when Porb ~4 minutes and forms thick disk

more massive WD accretes material from disk

model

helium ignites in shell at core-envelope boundary

helium-burning shell forces star to expand to yellow giant, ~103 yr

accretion turned off at selected final mass

0.4 Msun He-WD accretes He at 10-5 Msun/yr

shell burns inwards in series of mild flashes; lifts degeneracy

Helium core-burning star (sdB?) formed as shell reaches centre

V652 Her

V652 Her

~lo

g d

P/d

t

P dotlinear analysis of

evolutionary models gives fundamental pulsation period

dP/dt, derivative of period wrt time (or dP/dn) also obtained

evolution track through P-dP/dn diagram…

helium ignites in shell at core-envelope boundary

helium-burning shell forces star to expand to yellow giant, ~103 yr

accretion turned off at selected final mass

0.6 M CO-WD accretes He at 10-5 M /yr

0.5 M CO-WD

0.6 M , X=0.001

CO + CO WD merger

CO+He merger model: Test 2luminosity

EHe stars

CO+He mergers

solid: 0.6MCO+Hedashed: 0.5MCO+He

light: accretionheavy: contraction

EHes

Baade radii from pulsating EHes

HD168476

LS IV-1 2

Photospheric Abundances

Since 1995, analyses of optical and/or UV spectra to give abundances for 17 extreme helium stars, Pandey et al. (2006)

a) Proxies for metallicity (Ni,Mn,Cr,Fe)b) Overabundant light elements (Mg,Si,S,..?)

Photospheric Abundances II

c) C always overabundantd) N overabundant (=C+N+O)e) O ranges from overabundant to normal?ef) s-process “ “ “ “

CO+He merger model: Test 6abundances

 We consider the initial conditions and the mixing recipe adopted by Saio & Jeffery (2002; see

also Pandey et al. 2001). The atmosphere and envelope of the resultant EHe are composed of two zones from the accreted He white dwarf (He:H, He:He) and three zones from the C-O white dwarf (CO:H, CO:He, CO:CO)

Hydrogen —Surviving hydrogen is contributed by the layers He:H and CO:H.Helium —helium from the He:He layer effectively determines the final He abundance. Carbon — 12C is provided either by C from the top of the C-O white dwarf core or from carbon

in the CO:He layer as a result of He burningNitrogen.—The nitrogen (14N) is provided by the He:He and CO:He layers. This prediction,

which closely matches the observed N and He abundances at all Fe for all but three stars, requires almost complete conversion of initial C, N, and O to N, as assumed for the layer

He:He.

Oxygen.—The oxygen (16O) is assumed to be a product of He burning and to be contributed by either the CO : CO layer (recipe 1) or the CO : He layer (recipe 2). Where O is enriched, the O/C ratio is set by a simple ratio of mass fractions. The O abundance for most of the other EHes appears to be a star's initial abundance. This is hard to understand because observed N comes from destruction of C+O! What are implications of 18O in RCrB stars?

Yttrium and zirconium.—s-process enrichment is sited in the CO:He and CO:CO layers. An

enrichment by a factor of 10 in CO:He is insufficient to match observed Y and Zr without a greater enrichment in the CO:He layer or additional of material from the CO : CO layer. Significantly, the two most obviously s-process–enriched EHes are also among the most O-rich.

Pandey, Lambert, Jeffery & Rao 2006, ApJ 638, 454

not bad!

Observational tests for merged binary white dwarf models

He+He : V652 Her

mass

radius

luminosity pulsation period

dP/dt

composition

CO+He : EHe stars

absence of binaries radii and masses surface gravities contraction rates surface abundances number densities

angular momentum ?

Links with other classes

• N-rich

• Solid: C-rich ~1%

• Open: C-strong ~10%

Links with other classes

CO+He WD Merger RCrB EHe He-sdO+ O(He) ?? high-mass CO-WD (~0.8-0.9M)

• High-L : C-rich ~1%, O-rich: s-rich• High-L : C-rich ~1%, O-”normal”

He+He WD Merger ? EHe He-sdB He-sdO- CO-WD (~0.4-0.6M)

• Low-L: N-rich (V652 Her): more mass, more H• Low-L: C-rich (LSS 3184): less mass, less H• Result depends on WD progenitor masses and how much

H survives the merger

Hybrid H/He+He WD Merger ?

He+He WD merger

Conclusion

• Progress since 1995– Census of Teff, log g, abundances (a few mere facts!)

– Pulsation masses for several EHes– Quantitative predictions from WD merger models are very

promising– Improvements to models for atmospheres

• Challenges– Apply new models to the spectroscopy– Need better data and models for pulsation (masses)

– High Teff stars, winds

– Hot RCrBs– Establishing links to other H-def classes – FUSE :( WSO-UV :) GAIA :)– Merger model predictions of nucleosynthesis, mixing and

statistics

models for helium stars with CO coresSaio 1988

equilibrium models with degenerate CO core and helium envelope, allowed to contract as He-shell burns out

Mc-Ls relation for helium stars contraction rates for helium stars

models for helium stars with CO coresSaio 1988

equilibrium models with degenerate CO core and helium envelope, allowed to contract as He-shell burns out

contraction rates for helium stars

dTeff/dt K/yr 0.9Msun model HD160641 105 HD168476 50 BD-9 4395 108 BD+1 4381 30

sensitive test for more realistic models, if we could make them!

Hydrodynamic simulations

• Benz et al. 1990, ApJ 342, 986, & ApJ 348, 647

• Segretain et al. 1997, ApJ 481, 355: SPH - 0.9+0.6 M

t=27.2s

/109 cm

SPH Simulations: 0.8+0.6 T Isern & Guerrero 2002, WD13 Naples

Temporal evolution of accreting WD

convection zone

hydrogen-burning shell

helium-burning shell

Mi=0.6X=0.001

H ignition He ignition

IUE observations of helium stars

predicted contraction rates up to ~100 K/yr

150 observations of seventeen helium stars from 1979 to 1995

effective temperatures and angular diameters measured by fitting model atmospheres

vectors represent predicted temperature evolution over 10 years for 0.7 and 0.9 (solid) Msun helium stars respectively

dT/dt significant

Radius measurement (Baade’s method)

method

IUE SWP+LWP LORES fluxes + model atmospheres Teff and

integrating SAAO 1.9m radial velocities R

R=. R/ L = R2Teff

4

M= gR2/G

PV Tel + 2 others

CO+He merger model: Test 2 M, R and L

3 methods for estimating masses of EHes

Ms spectroscopic mass Mc-Ls & gMp pulsation mass & gMd direct mass R, , , g

Errors in masses are horrendous, because log g is involved.

Also get Teff, R and hence L. Errors much smaller.

Ms Mp Md

HD168476 0.95 0.85 0.94 BD+1 4381 1.09 0.93 0.07? LS IV-1 2 0.78 0.94 0.79

CO+He merger model: Test 2

EHe starsCO+He mergers

solid: 0.6MCO+Hedashed: 0.5MCO+He

light: accretionheavy: contraction

EHes

Baade radii from pulsating EHes

HD168476

LS IV-1 2

CO+He merger model: Test 3

gravitiesHD160641

HD168476

solid: 0.6 M CO+He

dashed: 0.5 M CO+He

dotted: 0.7 M He+He

CO+He merger model: Test 4

contraction

HD160641

BD-9 4395

HD168476

BD-1 3438

Can mergers really happen?Stable or Unstable Mass Transfer ?

If q > ~0.6, the increase of stellar radius exceeds the increase of Roche radius.

M-R relation for WD is due to requirement of hydrostatic balance, hence radius will increase on a dynamical timescale (~s), leading to runaway mass transfer. SPH calculations by Benz et al. (1990) show formation of thick disk around the more massive WD.

If q < ~0.6, stable mass transfer may produce an AM CVn type system

Conservation of Angular Momentum ?

When the surface velocity is close to the Kepler velocity, angular momentum is transported efficiently from the star to the disk so that accretion continues as long as matter around the star exists (Paczynski 1990, Pophan & Naryan 1990)

Angular Momentum

• Gourgouliatos & Jeffery (2006, submitted)• Compute specific angular momentum distribution

(SAM) of CO+He WDs at contact• Conserve SAM and compute distribution of mass in

disk• Accrete mass from disk to CO WD and compute

SAM• Compute surface rotation velocity for product as

– A) a giant– B) contraction to the WD sequence

• Compare with observed v sin i for EHes

• Repeat for He+He WD mergers