Extreme helium stars: a panorama Simon Jeffery - Armagh Observatory.
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Transcript of Extreme helium stars: a panorama Simon Jeffery - Armagh Observatory.
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