a-type stars: evolution, rotation and binarity
DESCRIPTION
A-type stars: evolution, rotation and binarity. Arlette Noels, Josefina Montalban Institut d’ Astrophysique et Géophysique Université de Liège, Belgium and. Carla Maceroni INAF - Rome Astronomical Observatory, Italy. THE A STAR PUZZLE - IAU Symposium 224 Poprad, Slovakia July 8, 2004. - PowerPoint PPT PresentationTRANSCRIPT
Arlette Noels, Josefina MontalbanArlette Noels, Josefina MontalbanInstitut d’ Astrophysique et Géophysique Université de Liège, BelgiumInstitut d’ Astrophysique et Géophysique Université de Liège, Belgium
andand
Arlette Noels, Josefina MontalbanArlette Noels, Josefina MontalbanInstitut d’ Astrophysique et Géophysique Université de Liège, BelgiumInstitut d’ Astrophysique et Géophysique Université de Liège, Belgium
andand
Carla MaceroniCarla MaceroniINAF - Rome Astronomical Observatory, ItalyINAF - Rome Astronomical Observatory, Italy
Carla MaceroniCarla MaceroniINAF - Rome Astronomical Observatory, ItalyINAF - Rome Astronomical Observatory, Italy
THE A STAR PUZZLE - IAU Symposium 224 Poprad, Slovakia July 8, 2004THE A STAR PUZZLE - IAU Symposium 224 Poprad, Slovakia July 8, 2004THE A STAR PUZZLE - IAU Symposium 224 Poprad, Slovakia July 8, 2004THE A STAR PUZZLE - IAU Symposium 224 Poprad, Slovakia July 8, 2004
Mass ~ 1.5 – 3 M
Teff ~ 7000 – 11000 KB-V ~ 0.0 – 0.30L ~ 10 – 50 L
Fundamental Fundamental parametersparametersFundamental Fundamental parametersparameters
age = 3.12 108 yr
age = 8.107 yr
age = 3.108 yr
H burning phaseH burning phaseH burning phaseH burning phase
Convective coreConvective core Convective coreConvective core X profile
Convective core: Convective core: temperature profiletemperature profile
Convective core: Convective core: temperature profiletemperature profile
Convective coreConvective core Convective coreConvective core
Overshooting Overshooting Overshooting Overshooting
t = 4. 106
t = 2. 107
H = 3.0 108 yrH = 2.2 108 yr
H = 2.9 108 yr
Overshooting Overshooting Overshooting Overshooting
Needed to fit CMD for open clusters and eclipsing binariesNeeded to fit CMD for open clusters and eclipsing binaries
Increases with mass (Andersen et al. 1990)Increases with mass (Andersen et al. 1990)
Overshooting Overshooting Overshooting Overshooting
NoNo isothermal core isothermal coreNoNo isothermal core isothermal core
Convective core: Convective core: temperature profiletemperature profile
Convective core: Convective core: temperature profiletemperature profile
Isothermal coreIsothermal core
Overshooting Overshooting Overshooting Overshooting
same size of He coresame size of He core same size of He coresame size of He core
Pre-main Pre-main sequencesequencePre-main Pre-main sequencesequence 1.5 – 4 M
Fully convectiveFully convective
Fully radiativeFully radiative
FormicolaFormicola et al. 2004et al. 2004
Pre-main Pre-main sequencesequencePre-main Pre-main sequencesequence
Palla & Stahler 1993Palla & Stahler 1993dM/dt = 10dM/dt = 10-5-5
Behrend & Maeder 2001,Behrend & Maeder 2001,dM/dt=1/3 (dM/dt)dM/dt=1/3 (dM/dt)discdisc
BirthlinesBirthlinesBirthlinesBirthlines
Pre-main Pre-main sequencesequencePre-main Pre-main sequencesequence
Effect of Effect of treatment of treatment of convection on convection on PMS evolutionary PMS evolutionary tracks location tracks location
Effect of Effect of treatment of treatment of convection on convection on PMS evolutionary PMS evolutionary tracks location tracks location
FST (Canuto et al. 1996).FST (Canuto et al. 1996).
MLT, MLT, =1.6=1.6MLT, MLT, =1.6=1.6
Convective Convective envelopeenvelopeConvective Convective envelopeenvelope
1.8 M
Convection in A-type star envelopes is superadiabaticConvection in A-type star envelopes is superadiabatic
> > > > > > > >
HI, HeIHI, HeIHI, HeIHI, HeI HeIIHeIIHeIIHeII
Thickness of the mixed layersThickness of the mixed layers Abundance anomaliesAbundance anomalies
Gravitational Gravitational settlingsettlingGravitational Gravitational settlingsettling
Microscopic Microscopic diffusiondiffusionMicroscopic Microscopic diffusiondiffusion
Radiative forcesadiative forces (Michaud et al. 1976, …) Turbulent transporturbulent transport (Schatzman 1969, Vauclair et al. 1978)
Radiative forcesadiative forces (Michaud et al. 1976, …) Turbulent transporturbulent transport (Schatzman 1969, Vauclair et al. 1978)
Enough but not too muchEnough but not too much Enough but not too muchEnough but not too much
Changes in the surface abundances (Richer et al. 2000)Changes in the surface abundances (Richer et al. 2000)
1.5M 1.7M2.5M
Changes in the internal structure
1. Mass of the convective envelope2. Fe convection zone around 200000 K
Changes in the internal structure
1. Mass of the convective envelope2. Fe convection zone around 200000 K
Rotation Rotation Rotation Rotation
A-type stars are rapid rotators: A-type stars are rapid rotators: vvrotrot up to 300 up to 300
km/s.km/s. Am and Ap: vAm and Ap: vrot rot < 120 km/s< 120 km/s
Normal A0-F0 stars : vNormal A0-F0 stars : vrot rot > 120 km/s (Abt & Morrel 1995)> 120 km/s (Abt & Morrel 1995)
A-type stars are rapid rotators: A-type stars are rapid rotators: vvrotrot up to 300 up to 300
km/s.km/s. Am and Ap: vAm and Ap: vrot rot < 120 km/s< 120 km/s
Normal A0-F0 stars : vNormal A0-F0 stars : vrot rot > 120 km/s (Abt & Morrel 1995)> 120 km/s (Abt & Morrel 1995)
Rotation Rotation Rotation Rotation
A-type stars are rapid rotators: A-type stars are rapid rotators: vvrotrot up to 300 up to 300
km/s.km/s. Am and Ap: vAm and Ap: vrot rot < 120 km/s< 120 km/s
Normal A0-F0 stars : vNormal A0-F0 stars : vrot rot > 120 km/s (Abt & Morrel 1995)> 120 km/s (Abt & Morrel 1995)
A-type stars are rapid rotators: A-type stars are rapid rotators: vvrotrot up to 300 up to 300
km/s.km/s. Am and Ap: vAm and Ap: vrot rot < 120 km/s< 120 km/s
Normal A0-F0 stars : vNormal A0-F0 stars : vrot rot > 120 km/s (Abt & Morrel 1995)> 120 km/s (Abt & Morrel 1995)Abt & Morrell 1995, Abt 1995:Abt & Morrell 1995, Abt 1995:Rotation alone can explain the occurrence of abnormal or Rotation alone can explain the occurrence of abnormal or normal normal main-sequence A stars because of our inability to main-sequence A stars because of our inability to distinguish marginal Am stars from normal ones in A2-F0 distinguish marginal Am stars from normal ones in A2-F0 and our inability to disentangle evolutionary effectsand our inability to disentangle evolutionary effects
Abt & Morrell 1995, Abt 1995:Abt & Morrell 1995, Abt 1995:Rotation alone can explain the occurrence of abnormal or Rotation alone can explain the occurrence of abnormal or normal normal main-sequence A stars because of our inability to main-sequence A stars because of our inability to distinguish marginal Am stars from normal ones in A2-F0 distinguish marginal Am stars from normal ones in A2-F0 and our inability to disentangle evolutionary effectsand our inability to disentangle evolutionary effects
BUTBUTBUTBUT
Debernardi & North 2001 V392 Carinae: vsini ~ 27 km/s no peculiarDebernardi & North 2001 V392 Carinae: vsini ~ 27 km/s no peculiarDebernardi & North 2001 V392 Carinae: vsini ~ 27 km/s no peculiarDebernardi & North 2001 V392 Carinae: vsini ~ 27 km/s no peculiar
Rotation Rotation Rotation Rotation
New Catalogue by Royer et al. 2002New Catalogue by Royer et al. 2002:
M > 1.6MM > 1.6M or B-V < 0.25-0.3: or B-V < 0.25-0.3: Little or no stellar activityLittle or no stellar activity No evidence of significant angular momentum lossNo evidence of significant angular momentum loss There is no trend on rotation with age (vsin i ~ cte) There is no trend on rotation with age (vsin i ~ cte)
M < 1.6MM < 1.6M or B-V > 0.25-0.3: or B-V > 0.25-0.3: Stellar activity does not depend on age or rotationStellar activity does not depend on age or rotation Very slow angular momentum loss. Braking time ~ 10Very slow angular momentum loss. Braking time ~ 1099yryr.
M > 1.6MM > 1.6M or B-V < 0.25-0.3: or B-V < 0.25-0.3: Little or no stellar activityLittle or no stellar activity No evidence of significant angular momentum lossNo evidence of significant angular momentum loss There is no trend on rotation with age (vsin i ~ cte) There is no trend on rotation with age (vsin i ~ cte)
M < 1.6MM < 1.6M or B-V > 0.25-0.3: or B-V > 0.25-0.3: Stellar activity does not depend on age or rotationStellar activity does not depend on age or rotation Very slow angular momentum loss. Braking time ~ 10Very slow angular momentum loss. Braking time ~ 1099yryr.
Rotation on Rotation on MSMSRotation on Rotation on MSMS
Rotational velocity distribution Rotational velocity distribution must be imposed the pre-main sequence evolutionmust be imposed the pre-main sequence evolution
(Wolff & Simon 1997)
Rotation in PMSRotation in PMSRotation in PMSRotation in PMS
Importance of the Birthline locationImportance of the Birthline locationImportance of the Birthline locationImportance of the Birthline location
From vsini in 145 From vsini in 145 in Orion (1 Myr), Wolff et al. 2004in Orion (1 Myr), Wolff et al. 2004:
1.1. Braking of stars with M< 2 MBraking of stars with M< 2 M as they evolve down as they evolve down their convective tracks (disk interaction)their convective tracks (disk interaction)
2.2. Conservation of angular momentum as stars evolve Conservation of angular momentum as stars evolve long their radiative traks long their radiative traks
1.1. Braking of stars with M< 2 MBraking of stars with M< 2 M as they evolve down as they evolve down their convective tracks (disk interaction)their convective tracks (disk interaction)
2.2. Conservation of angular momentum as stars evolve Conservation of angular momentum as stars evolve long their radiative traks long their radiative traks
High accretion rate birthline High accretion rate birthline at larger R at larger RLow accretion rate birthline Low accretion rate birthline at radiatively low R at radiatively low R
Rotation: effect on stellar evolutionRotation: effect on stellar evolutionRotation: effect on stellar evolutionRotation: effect on stellar evolution
Surface effects:Surface effects: Photometric parameters Photometric parameters Anisotropic mass lossAnisotropic mass loss
Departure from sphericity: meridional circulationDeparture from sphericity: meridional circulation Differential rotation and instabilities Differential rotation and instabilities (e.g. Pinsonneault (e.g. Pinsonneault
1997)1997) Transport of angular momentum and chemicalsTransport of angular momentum and chemicals
Similar to overshooting in the HRDSimilar to overshooting in the HRDButBut
Different internal structure?Different internal structure?
Maeder & Zahn (1998), Zahn (1992) Maeder & Zahn (1998), Zahn (1992)
Rotation: effect on stellar evolutionRotation: effect on stellar evolutionRotation: effect on stellar evolutionRotation: effect on stellar evolution
Palacios et al. 2003
2.2 M2.2 M
1.8 M1.8 M
1.5 M1.5 M
1.4 M1.4 M
1.35M1.35M
Time spent on MS increases by 20% in lower mass stars 10% in higher mass models
Transport by meridional circulationTransport by meridional circulationand highly anisotropic turbulence and highly anisotropic turbulence in a rotating and non magnetic star.in a rotating and non magnetic star.
Transport by meridional circulationTransport by meridional circulationand highly anisotropic turbulence and highly anisotropic turbulence in a rotating and non magnetic star.in a rotating and non magnetic star.
““New prescription of DhNew prescription of Dhkeeps the size of the keeps the size of the core”core” (Maeder 2003)
Maeder & Zahn 1998Maeder & Zahn 1998Maeder & Zahn 1998Maeder & Zahn 1998
Maeder 2003Maeder 2003Maeder 2003Maeder 2003 No rotationNo rotationNo rotationNo rotation
Rotation: effect on stellar evolutionRotation: effect on stellar evolutionRotation: effect on stellar evolutionRotation: effect on stellar evolution
Dh Dh Maeder 2003Maeder 2003 >> Dh >> Dh Maeder & ZahnMaeder & Zahn Dh Dh Maeder 2003Maeder 2003 >> Dh >> Dh Maeder & ZahnMaeder & Zahn
Maeder (2003): Maeder (2003): balance balance between horizontal turbulence between horizontal turbulence and excess of energy in the and excess of energy in the differential rotation differential rotation
Maeder (2003): Maeder (2003): balance balance between horizontal turbulence between horizontal turbulence and excess of energy in the and excess of energy in the differential rotation differential rotation
Horizontal turbulent Horizontal turbulent diffusivity: Dhdiffusivity: DhHorizontal turbulent Horizontal turbulent diffusivity: Dhdiffusivity: Dh
Vertical effective diffusivity: DeffVertical effective diffusivity: DeffVertical effective diffusivity: DeffVertical effective diffusivity: Deff
Rotation: effect on stellar evolutionRotation: effect on stellar evolutionRotation: effect on stellar evolutionRotation: effect on stellar evolution
Mathis & Zahn 2004Mathis et al. 2004
ββ-viscosity prescription to determine Dh-viscosity prescription to determine Dhββ-viscosity prescription to determine Dh-viscosity prescription to determine Dh
Differential rotation in radiative layers Differential rotation in radiative layers (Tayler instability)(Tayler instability) Magnetic field Magnetic field (Spruit 1999, 2002).(Spruit 1999, 2002).
Magneto-rotational instability Magneto-rotational instability (Balbus & Hawley 1991)(Balbus & Hawley 1991) could could transport J to the surface transport J to the surface (Arlt et al. 2003).(Arlt et al. 2003). Timescale ~ life time for A type stars Timescale ~ life time for A type stars Effect on J of Ap starsEffect on J of Ap stars
Rotation: effect on stellar evolutionRotation: effect on stellar evolutionRotation: effect on stellar evolutionRotation: effect on stellar evolution
Interaction rotation-convectionInteraction rotation-convectionInteraction rotation-convectionInteraction rotation-convection
Convective envelope:Convective envelope:
Reduce the size of the overshooting layer at the Reduce the size of the overshooting layer at the bottom of the convective envelope bottom of the convective envelope (Chan 1996, Julien et al. 1996)(Chan 1996, Julien et al. 1996)
Convective core Convective core ((Browning et al. 2004):Browning et al. 2004):
Differential rotationDifferential rotation Overshooting Overshooting
Rotation: open Rotation: open questionsquestionsRotation: open Rotation: open questionsquestions
Overshooting and/or rotatonal mixing in theOvershooting and/or rotatonal mixing in the
internal regions?internal regions? Mixing close to the surface:Mixing close to the surface:
Li, Be in A-type stars and in the SunLi, Be in A-type stars and in the Sun Am surface abundances (D ~ Am surface abundances (D ~
D(He)D(He)00((//oo))nn))
Transport of angular momentum in the Transport of angular momentum in the
radiative radiative
regions internal rotation in A-type stars: regions internal rotation in A-type stars: solid or differential rotation?solid or differential rotation? role of magnetic instabilities role of magnetic instabilities
Puzzle pieces (general trends)Puzzle pieces (general trends)
A Am Ap
(Sr-Cr,Si) Ap (HgMn)
close close binary binary
frequencyfrequency normnorm
VeryVery
highhighLowLow NormNorm
rotationrotation FastFast SlowSlow SlowSlow SlowSlow
magneticmagnetic
fieldsfields nono nono yes, strongyes, strong nono
Binarity slowing-down of rotation Am phenomenon
magnetic Ap’s: strong magnetic fields binarity
Typically (~ not far from always) Am’s are (close) binaries Rarely Ap’s are binaries, and anyway with an orbital P≥2.5d
Questions on binarity:is binarity a necessary and sufficient condition to be an Am ?
is binarity - through syncronization and circularization mechanisms - just an efficient brake of stellar rotation or does it affect the stellar structure in other ways?
A-type star binarity/non-binarityA-type star binarity/non-binarity
Perhaps
..no definite answer…
...l king for the answers
The synchronization (and circularization) theories are usually compared withthe Observed (orbital) Period Distributions (OPD), the rotational data and the eccentricity - P plots.
Three sorts of problems:Limits of the available theories or in their
applicationSmall and non homogeneous available samples
with sufficiently accurate elementsSelection effects on the OPD
In late-type stars it is the turbulent dissipationturbulent dissipation in the outer convection zone that retards the equilibrium tide,
Synchronization & circularization Synchronization & circularization theories: theories:
I. Zahn’s tidal mechanismsI. Zahn’s tidal mechanisms
Synchronization & circularization Synchronization & circularization theories: theories:
I. Zahn’s tidal mechanismsI. Zahn’s tidal mechanisms
In early type stars the dissipation mechanism is radiative radiative dampingdamping, which acts on the dynamical tide (forced gravity waves are emitted from a lagging convective core and damped in the outer layers).
Two necessary ingredients:
tidal bulges
dissipation mechanism
non-alignement torque
a
R
Ω
ω
,611
6222
a
R
I
MRq
t
k
dt
d
t fsync
,151 2
17
22
652
21
3
a
RE
I
MRqq
R
GM
tsync
,12
211 221
22
6112
21
3
a
RE
I
MRqq
R
GM
tcirc
8
2 1 2
21eln 1
a
Rqq
t
k
dt
d
t fcirc
Zahn tidal theory: timescales
E2 is a constant strongly dependent on the size of the convective dependent on the size of the convective corecore
time friction typical the and constant apsidal 2 the is , ndf tkmmq 221
related to the density profilerelated to the density profile inside the star
In early type stars the timescales increase more rapidly In early type stars the timescales increase more rapidly with a (or P) and the forces have a shorter rangewith a (or P) and the forces have a shorter range
Late-type stars:
Early –type stars:
,a
R
R
MLq)q ( cost
tγ
sync
N8
3381
83
4
9
2
1101
II. Tassoul’s hydrodynamical theoryII. Tassoul’s hydrodynamical theory
Transient strong meridional currents, produced by the tidal action, transfer angular momentum between the stellar interior and the Ekman layer close to the surface. If ω>Ω the star spins down.
Timescales:
8498
1
811
4
9
22)1(10
1
a
R
R
MLβq cost'
tγ
circ
N
,logr
vN
where is the eddy and the radiative viscosity of the outer layers (N=0 for radiative envelopes).
rv
with
γ takes somehow into account the fact that the eqs are solved for ~circular and ~synchronized motions.
Tassoul’s mechanism has a longer range and a much higher Tassoul’s mechanism has a longer range and a much higher efficiency for early-type starsefficiency for early-type stars
Warnings!
! the use of timescales cannot replace the integration of the evolutionary equations, which require as well the introduction of stellar evolutionary models (see Claret et al. 1995, Claret et Cunha 1997)
! Both theories are for quasi-circular & quasi-synchronized orbits. Tassoul introduces an arbitrary factor (~10-40) in the timescales.
! The strong dependence of the processes on R/a requires systems with very accurate element determination.
log (t/tcri) log (t/tcri)
e e
Application to A and early type stars(Matthews & Mathieu 1992, Claret et al. 1995, 1997)
Zahn Tassoul, = 1.6
t: binary age, tcrit : time for circularization.
Circ.Non-circ.
From Claret et al. 1995, 1997
Application to A and early type stars, II(Matthews & Mathieu 1992, Claret et al. 1995, 1997)
Tassoul, = 1.6Tassoul, = 0
log (t/tcri) log (t/tcri)
e e
Spin – orbit synchronization:Am
orbsyn P
Riviv 6.50)90(sin
(Am sample from Budaj 96 (Segewiss 93) + Updated v sin i from Royer et al. 2002)
Expected syncronization P: R/a≈0.25 ( North & Zahn 02)
M=2.0 R=3.0 q=0.2
R=2.1 q=1.0
ω=ΩIn a synchronized binary:
d 21
23
125.0
1159.0 1
qmR
Psyn
Spin – orbit synchronization Am (Am sample from Budaj 96 (Segewiss 93) + Updated v sin i from Royer et al. 2002)
Expected syncronization P: R/a≈0.25 ( North & Zahn 02)
M=2.0 R=3.0 q=0.2
R=2.1 q=1.0
ω=Ω
v sin i before updating
Empty region
P-dependent tidal mixing
Spin – orbit synchronization Am (Am sample from Budaj 96 (Segewiss 93). Updated v sin i from Royer et al. 2002)
Expected syncronization P: R/a≈0.25 ( North & Zahn 02)
M=2.0 R=3.0 q=0.2
R=2.1 q=1.0
ω=Ω
Selection effects on SB’sSelection effects on SB’s
days 31
22
3316
1max)1()1(
sin1063.9)(
23
Kqe
iqmKP
minimum observable radial velocity amplitude, K1≠ instr. limit
maximum observable orbital Period:
P=P(m1,q,e) [ sin i =1.0]
if K1 =10 Km/s
detailed modeling of SB8 selection effects (Hogeveen 1992) suggests for A-type stars:
K1≈ 25 Km/s
SB1 q distribution is peaked around q≈0.2.
m1=2.0
Missed SB1