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White Dwarfs With contributions from S. R. Kulkarni T. Monroe.

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Presentation on theme: "White Dwarfs With contributions from S. R. Kulkarni T. Monroe."— Presentation transcript:

1 White Dwarfs With contributions from S. R. Kulkarni T. Monroe

2 References D. Koester, A&A Review (2002) “White Dwarfs: Recent Developments” Hansen & Liebert, Ann Rev A&A (2003) “Cool White Dwarfs” Wesemael et al. PASP (1993) “An Atlas of Optical Spectra of White- Dwarf Stars” Wickramsinghe & Ferrario PASP (2000) “Magnetism in Isolated & Binary White Dwarfs”

3 References Dreizler, S. 1999, RvMA, 12, 255D Fontaine et al. 2001, PASP, 113, 409 Hansen, B. 2004, Physics Reports, 399, 1 Hansen, B & Liebert, J. 2003 ARA&A, 41, 465 Hearnshaw, J.B. 1986, The Analysis of Starlight. Koester, D. & Chanmugam, G. 1990, RPPh, 53, 837K Shipman, H. 1997, White Dwarfs, p. 165. Kluwer Wesemael et al. 1993, PASP, 105, 761

4 How stars die Stars above 8 Msun form neutron stars and black holes Below 8 Msun the stars condense to O-Ne-Mg white dwarfs (high mass stars) or usually C-O white dwarfs Single stars do not form He white dwarfs but can form in binary stars We know of no channel to form H white dwarfs of some reasonable mass

5 History of White Dwarf Discovery Bessell (1844)-variability in proper motions of Sirius and Procyon  dark companions Clark (1861) visually sighted Sirius B Schaeberle (1896) Lick Obs. announced Procyon’s companion 40 Eri (faint white and red stars) –Class A0, Russell dismissed when 1 st Russell diagram published –Adams confirmed A-type Adams (1915)-Sirius B spectrum  Type A0 Eddington (1924) Mass-Luminosity Relationship –Coined “white dwarfs” for 1 st time –Deduced mass and radius of Sirius B  density=53,000x water Fowler (1926) WDs supported by electron degeneracy pressure, not thermal gas pressure Chandrasekhar (early 1930s) worked out details of white dwarf structure, predicted upper mass limit of 1.44 M sun, & found mass-radius relation

6 Early Classifications Kuiper (mid-1930s, Lick Obs.) WDs found in increasing numbers –1941 introduced 1 st WD classification scheme w in front of spectral type and Con stars Luyten (1921) proper motion studies from faint blue star surveys –1952 presented new scheme for 44 WDs D for true degeneracy, followed by A, B, C, or F Greenstein (1958) introduced new scheme –9 types

7 Current Classifications Sion (et al. 1983) ~2200 WDs w/in ~500 pc of Sun D=degenerate Second Letter-primary spectroscopic signature in optical –DA-Hydrogen lines (5000K45,000K) Additional letters indicate increasingly weaker or secondary features, e.g. DAZ, DQAB –P-polarized magnetic, H-non-polarized magnetic, V-variable T eff indicated by digit at end; 50,400/T eff, e.g. DA4.5 New class T eff <4000K, IR absorption for CIA by H 2

8 DA Spectra Rapid settling of elements heavier than H in high gravity DB Spectra

9 DQ Stars & Spectra Helium-rich stars, generally characterized by C 2 -Swan bands Hotter DQs have C I

10 PG 1159 Spectra Features due to CNO ions, T eff >100,000K Absence of H or He I features; He II, C IV, O VI ZZ Ceti

11 Magnetic WDs About 5% of field white dwarfs display strong magnetism 3 classes of H- atmosphere MWDs based on field strength He-atmosphere MWDs have unique features

12 Basic Picture 75% DA, 25% non-DA Spectral classification provides info about principal constituent, with some T info Progenitors: Post-AGB stars, central stars of planetary nebulae (CSPN), hot subdwarfs Expected structure-stratified object with ~0.6M sun –C-O core, He-rich envelope, H-rich shell O-Ne cores-most massive –Atmosphere contains <10 -14 M Many WDs have pure H or He atmospheres Thicknesses of H and He

13 Mechanisms in Atmosphere Gravitational diffusion Convection Radiative levitation Magnetism Accretion Wind-loss T-sensitive  T determines chemical abundances

14 Effects of Mechanisms Diffusion & Settling –Gravitational separation leads to pure envelope of lightest element t<10 8 yr But, observations show traces of heavier elements –radiative levitation –Cooler WDs result of recent accretion event Radiative Levitation T>40kK –Radiative acceleration on heavy elements Convection for T<12kK –Convection zone forms and increases inward as star cools –For He envelopes, convection begins at high T –Mixing changes surface composition –Need to couple models of atmospheres and interiors

15 Statistics T>45kK DA far outnumber DO –Ratio increases to about 30kK (diffusion) DB gap in 45k-30kK range –Float up of H Always enough H to form atmosphere? –Dredge up of He T<30kK He convection zone massive engulfs outer H layer if thin –30kK-12kK 25% stars revert to DB spectral type (edge of ZZ Ceti Strip) –Convection zone increases as T decreases. At T~11kK, numbers of DAs and non-DAs are ~equal (ZZ Ceti Strip) ‘Non-DA gap’ for 5000-6000K dearth of He atmospheres

16 Spectral Evolution Gaps  individual WDs undergo spectral evolution –Compositions change, DA  DB  DA, as T changes Evolution of convection zone? Accretion? Explanation of ‘non-DA gap’-opacity? Bergeron et al. –Low opacity of He I means small amounts of H dominates opacity –H - atomic energy levels destroyed when H added to dense atmosphere-reduces H opacity contribution –Must accrete a lot of H to make difference in photospheric conditions  DA (fixes 6000K edge) –Re-appearance of DBs at 5000K b/c convection zone grows, H is diluted with additional He –This fails! Destruction of H - bound level produces free e -, which provide opacity

17 ZZ Ceti Cooling EvolutionCSPN Hot DAZs (T>40kK) Radiative leviation makes Z No Z cooler than 35kK ZZ Ceti w/ variable H layers 10 -8 …………………10 -4 Msun He-Rich DA (0.01 { "@context": "http://schema.org", "@type": "ImageObject", "contentUrl": "http://images.slideplayer.com/10/2873819/slides/slide_17.jpg", "name": "ZZ Ceti Cooling EvolutionCSPN Hot DAZs (T>40kK) Radiative leviation makes Z No Z cooler than 35kK ZZ Ceti w/ variable H layers 10 -8 …………………10 -4 Msun He-Rich DA (0.0140kK) Radiative leviation makes Z No Z cooler than 35kK ZZ Ceti w/ variable H layers 10 -8 …………………10 -4 Msun He-Rich DA (0.01

18 Model Atmospheres Plane-parallel geometry Hydrostatic equilibrium (mass loss rates) NLTE Stratisfied Atmospheres –Parameters: degree of ionization, intensity of radiation field Make radiative cross sections of each element depth dependent Convection –Parameters of Mixing Length theory

19 White Dwarfs in Globular Clusters

20 Cluster White Dwarf Spectroscopy

21 White Dwarfs in Clusters Chronometers: Use cooling models to derive the ages of globular clusters Yardsticks: Compare nearby and cluster white dwarfs. Forensics: Diagnose the long dead population of massive stars

22 The Globular Cluster M4 Fainter white dwarfs are seen in this nearby cluster -> age = 12.7 +/- 0.7 Gyr M4 formed at about z=6 Disk formed at about z=1.5 dN/dM, differential mass spectrum dN/dM propto M -0.9

23 White Dwarfs in Open Clusters Open Clusters have a wide range of ages (100 Myr to 9 Gyr, the age of the disk) Use white dwarfs as chronometers Derive initial-mass to final-mass mapping Key Result: M WD about 8 M Sun This result is in agreement with stellar models

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26 Field White Dwarfs Identified by large proper motion yet faint object LHS (Luyten Half Second) NLTT (New Luyten Two Tenths) Blue Objects (found in quasar surveys) Very Hot objects (found in X-ray surveys)

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28 Field White Dwarfs

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30 Old White Dwarfs Microlensing observations indicate presence of 0.5 Msun objects in the halo Old white white dwarfs expected in our disk, thick disk and halo These old white dwarfs are paradoxically blue (cf cool brown dwarfs)

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33 Determination of Mass (Field Objects) Spectroscopic Method: Line (Hydrogen) width is sensitive to pressure which is proportional to gravity g = GM/R 2 Photometric Method: Broad-band photometry fitted to black body yields Teff and angular size Combine with parallax to get radius R Use Mass-Radius relation to derive Mass

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37 Masses of White Dwarfs

38 Magnetism in Isolated White Dwarfs About 5% of field white dwarfs exhibit strong magnetism On average, these white dwarfs have larger mass Some rotate rapidly and some not at all Magnetism thus influences the initial-final mapping relation Or speculatively, some of these are the result of coalescence of white dwarfs

39 Zeeman (Landau) Splitting

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41 Future/Active Work Exact masses of H and He layers –Thin or Thick Envelopes Explanations for DB-gap Explanations for ‘non-DA gap’ DAs outnumber He-rich WDs, yet progenitor PNN have ~equal numbers of H- and He-rich stars. What rids degenerates of He? Couple core & atmosphere models


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