evolution and internal structure of red giants
DESCRIPTION
EVOLUTION AND INTERNAL STRUCTURE OF RED GIANTS. Maurizio Salaris Astrophysics Research Institute Liverpool John Moores University. What stars am I going to talk about?. Kallinger et al. (2010). Adapted from Gallart (1999). From McConnachie et al. (2006). RGB stars as metallicity indicators. - PowerPoint PPT PresentationTRANSCRIPT
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EVOLUTION AND INTERNAL STRUCTURE OF RED
GIANTS
Maurizio SalarisAstrophysics Research InstituteLiverpool John Moores University
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Kallinger et al. (2010)
What stars am I going to talk about?
Adapted from Gallart (1999)
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From McConnachie et al. (2006)
RGB stars as metallicity indicators
↓
RC stars as SFH diagnostics
←
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Number of RGB stars above the HB level
reduced by a factor 2
10
14Gyr
[Fe/H]= −0.7 −0.35 +0.06 +0.40
Blue circles Teff + 100 K (only RGB)
8
6
3
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OUTLINE
• The three classes of red giant stars• Behaviour in the CMD (or HRD)• Internal structure• Some long-standing uncertainties• Examples of applications to more
general astrophysical problems
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Solar initial chemical composition
1.0 Mo
2.4 Mo
RGB
RC
EAGB
Representative evolutionary tracks
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RGB stars
• Objects with (initial) mass lower than ~2.0Mo
• Electron degenerate (nearly) isothermal He-core surrounded by a thin (~0.001-0.0001 Mo thickness) H-burning shell that is, in turn, surrounded by an extended convective envelope
• Evolution towards increasing luminosity and moderately decreasing Teff due to the steady increase of the He-core mass
• Efficient mass loss from the convective envelope• He-flash terminates RGB evolution when Mc
He~0.47 - 0.50Mo
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LIFETIMES
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RGB in the CMD (or HRD) metallicity age initial helium
One has to be careful with the intermediate ages
LLbol bol of the TRGB of the TRGB increases with increases with increasing Z, but increasing Z, but the behaviour in the behaviour in the CMD the CMD depends on the depends on the passbandpassband
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First dredge-up
After the 1st dredge up
12C/13C to ~ 25 from ~90
14N by a factor ~ 2 12C by ~ 30
%
7Li by a factor ≈ 20
16O
Y by 0.01 0.02
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RGB bump
1Mo solar composition
The size of the H-abundance discontinuity determines the ‘area’ of the bump region in the LF.
The shape of the H-profile discontinuity affects the shape of the bump region
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13 Gyr
Z=0.008
Z=0.0004
10 – 13 Gyr
Z=0.008
Dependence of the bump luminosity on age and metallicity
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Zoccali et al. (1999)
RGB bump detection in stellar populations
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Chemical profiles and energy generation
grad(T)=grad(T)=ln(T)/ln(T)/ln(P)ln(P)
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He-core Superadiabatic region
surface
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Difficulties with the parametrization of the RGB mass loss
13RGB 4 10 R
dM L
dt gR
He WD He WD limitlimit
red red HBHB
EHB EHB limitlimit
10,00010,000 KK
RR RR LyraeLyrae
The Reimers’ law
free parameter
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From Castellani & Castellani (1993)
Only extreme values of η affect appreciably the HRD of RGB stars
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!
Different parametrizationsDifferent parametrizations
dMR
dt 3.2
“Goldberg formula”
“Mullan’s formula”
32
dM g
dt R
-0.9
“Judge & Stencel formula”
dMg
dt -1.6
Catelan (2009)
dM L
dt gR
1.4
“modified Reimersformula”
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Uncertain dependence Uncertain dependence on the metallicityon the metallicity
Origlia et al. 2007
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From Salaris et al (1993)
Difficulties with the Teff scale of RGB models
The Teff scale of RGB models depends on:
i) Low-T opacities
ii) Treatment of superadiabatic gradient
iii)Boundary conditions
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Superadiabatic convection: The mixing length theory (Böhm-Vitense 1958)
a b c α
BV58 ⅛ ½ 24 calibration
HVB65 ⅛ ½ modified
calibration
ML1 ⅛ ½ 24 1.0ML2 1 2 16 0.6 -
1.0
ML3 1 2 16 2.0 ML2 and ML3 increase the convective efficiency compared to ML1
ll==αα HHp mixing lengthp mixing length
Widely used in stellar evolution codesWidely used in stellar evolution codes
Simple, local, time independent model, that assumes convective Simple, local, time independent model, that assumes convective
elements with mean size l,elements with mean size l, of the order of their mean free pathof the order of their mean free path
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The value of The value of αα affects strongly the effective affects strongly the effective temperature of stars with convective temperature of stars with convective envelopes envelopes
The’canonical’ The’canonical’ calibration is based on calibration is based on reproducing the solar reproducing the solar radius with a radius with a theoretical solar theoretical solar models (Gough & models (Gough & Weiss 1976)Weiss 1976)
We should always keep We should always keep in mind that there is a in mind that there is a priori no reason why priori no reason why αα should stay constant should stay constant within a stellar within a stellar envelope, and when envelope, and when considering stars of considering stars of different masses different masses and/or at different and/or at different evolutionary stagesevolutionary stages
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Are different formulations of the MLT Are different formulations of the MLT equivalent ?equivalent ?
Gough & Weiss (1976), Pedersen Gough & Weiss (1976), Pedersen et al. (1991)et al. (1991)
A simple testA simple test
The mixing length The mixing length calibration preferred in calibration preferred in White Dwarf model White Dwarf model atmospheres and atmospheres and envelopes (e.g. Bergeron envelopes (e.g. Bergeron et al. 1995) is the ML2 et al. 1995) is the ML2 with with αα=0.6=0.6
Salaris & Cassisi (2008)
ML2, ML2, αα=0.63 (solid - solar calibration)=0.63 (solid - solar calibration)ML1, ML1, αα=2.01 (dashed – solar calibration=2.01 (dashed – solar calibration))
ML2 models at most ~50 K hotter
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Hydro-calibration Extended grid of 2D
hydro-models by Ludwig, Steffen & Freytag
Static envelope models based on the mixing length theory calibrate α by reproducing the entropy of the adiabatic layers below the superadiabatic region from the hydro-models.
A relationship α=f(Teff,g) is produced, to be employed in stellar evolution modelling
(Ludwig et al.1999)From Freytag & Salaris (1999)
Previous attempts by Deupree & Previous attempts by Deupree & Varner (1980) Lydon et al (1992, Varner (1980) Lydon et al (1992, 1993)1993)
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Calibration of the mixing Calibration of the mixing length parameter using length parameter using RGB starsRGB stars
Effective Effective temperaturestemperatures
Prone to uncertainties in Prone to uncertainties in the temperature scale, the temperature scale, metallicity scale, colour metallicity scale, colour transformationstransformations
ColoursColours
CALIBRATION OF THE MIXING LENGTH CALIBRATION OF THE MIXING LENGTH ON RGB STARSON RGB STARS
See Paolo talk for more
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Solar calibrated models with different boundary conditions predict different RGB temperatures
Salaris et al. (2002)Salaris et al. (2002)
Montalban et al. (2004)Montalban et al. (2004)
Boundary Boundary conditionsconditions
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Calculations with empirical solar T() and same opacities as in model atmosphere (solid line) , compared with the case of
boundary conditions from detailed model atmospheres (ATLAS 9 – dashed line)
From Pietrinferni et al. From Pietrinferni et al. (2004)(2004)
Boundary Boundary conditions conditions taken at taken at =56=56
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The need for additional element transport mechanisms
Mucciarelli, Salaris et al. (2010)
Gratton et al. (2000)
Field halo starsField halo stars
Globular cluster M4Globular cluster M4
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From Salaris, Cassisi & Weiss From Salaris, Cassisi & Weiss (2002)(2002)
0.8 M0.8 Moo metal poor RGB metal poor RGB model model
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“The H-burning front moves outward into the stable region, but preceding the H-burning region proper is a narrow region, usually thought unimportant, in which 3He burns.
The main reaction is 3He (3He, 2p)4He: two nuclei become three nuclei, and the mean mass per nucleus decreases from 3 to 2. Because the molecular weight (µ) is the mean mass per nucleus, but including also the much larger abundances of H and 4He that are already there and not taking part in this reaction, this leads to a small inversion in the µ gradient. “
Eggleton et al. (2006)
ADDITIONAL TRANSPORT MECHANISMS
1Mo solar composition
See Corinne talk for more details
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Surface abundance variations on the RGB for the model with diffusion (red line) and the model without diffusion (blue line).
From Michaud et al. (2007)
0.8Mo Z=0.0001
ATOMIC DIFFUSION
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Effect of smoothing the H-profile discontinuity Cassisi, Salaris & Bono
(2002)
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Needs more than ~ 120 stars within ±0.20 mag of the bump peak, and photometric errors not larger than 0.03 mag to reveal the effect of smoothing lengths ≥ 0.5 Hp
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Salaris et al. (2002)
Bellazzini et al. (2001)
8,10,12,14 Gyr
Z=0.0004
Z=0.0002 0.00810 Gyr
TRGB as distance indicator
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Holtzman et al. (1999)
RGB stars in composite stellar populations, an example
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Synthetic MI -(V–I) CMD detailing the upper part of the RGB, and two globular cluster isochrones for [Fe/H] equal to −1.5 and −0.9, respectively
Metallicity distribution of the synthetic upper RGB CMD.
Salaris & Girardi (2005)
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Z=0.019
Girardi (1999)Solid lines end when 70 of tHe is reached. Short-dashed lines denote the evolution from 70 up to 85 of tHe, whereas the dotted ones go from 85 to 99 of
tHe.
RED CLUMP STARSRC stars are objects in the central He-burning phase.
A convective He-burning core is surrounded by a H-burning shell.
Above the H-burning shell lies a convective envelope
The path in the HRD is determined by the relative contribution of the central and shell burning to the total energy output
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He core mass at He ignition
Salaris & Cassisi (2005)Salaris & Cassisi (2005)
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INSIDE A RC STAR
Log(L/Lo)=1.7
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COMPARISON WITH RGB STARS
Log(L/Lo)=1.7
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log
( clo
g( c 2
)2)-1
5-1
5
Comparison of sound speed profilesComparison of sound speed profiles
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Treatment of Core ConvectionTreatment of Core Convection
C produced by He-burningC produced by He-burning
Opacity increasesOpacity increases
Radiative gradient discontinuity atRadiative gradient discontinuity at
the convective core boundarythe convective core boundary
Mass of Mass of convective convective core core increasesincreases
See, e.g. Castellani et al. (1971)See, e.g. Castellani et al. (1971)
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Michaud et al. (2008) have shown that the phase of core expansion can be also produced by atomic diffusion
What happens next ?
See Achim talk
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Typical evolution of temperature gradients and He abundances in the core of RC stars
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The RC age-magnitude-colour distribution for a given SFH depends on the trend of the TO lifetime with mass, and the He-burning/TO lifetime ratio with mass
Solar neighbourhood RC simulation (Girardi & Salaris 2001)
INPUT (Rocha-Pinto et al. (2000)
OUTPUT
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DIFFERENT SFHs PRODUCE VERY DIFFERENT RC MORPHOLOGIES
Girardi & Salaris (2001)
Solar neighbourhood
LMC fields
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ΔMλRC=ΔMλ
RC (local)-ΔMλRC (pop)
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Early AGB
Early-AGB stars are objects with an electron degenerate CO-core embedded within the original He-core at He-ignition.
An H-burning shell is efficient above the He-core boundary, surrounded by a convective envelope.
The evolution is similar to RGB stars. The early-AGB ends with the ignition of the He-burning shell (AGB clump). Timescales ≈107 yr
EARLY-AGB
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Internal stratification
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Development of CO-core degeneracy
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Log(L/Lo)=2.1
Comparison with RGB stars
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log
( clo
g( c 2
)2)-1
5-1
5
Comparison of sound speed profilesComparison of sound speed profiles
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Open questions:
• Accuracy of model Teff (superadiabatic convection + boundary conditions)
• RGB mass loss• Element transport mechanisms during
the RGB• Mixing in the core during the central
He-burning phase
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The role of red giants in population synthesis
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Difference between the theoretical I(TRGB) for t= 12.5 Gyr, scaled solar [Fe/H]=−1.38 and the theoretical values predicted for the ages and scaled solar [Fe/H] values displayed.
The underlying theoretical models are from Girardi et al. (2000). Panels (a), (b), (c) and (d) show, respectively, the results using the Yale transformations, Westera et al. (2002) transformations, Girardi et al. (2002) transformations, and the transformations used in Girardi et al. (2000).
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Michaud et al. (2010)
ATOMIC DIFFUSION AND INTERNAL PROFILES
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The properties of H-burning shell, hence the luminosity of the RGB star are mainly determined by the mass (Mc
He) and radius (Rc
He) of the He-core.
Using the M-R relation of cold WDs, the CNO-cycle energy generation mechanism and electron scattering opacity in the shell, homology considerations give:
dln(L)/dln(McHe) ≈ 8 – 10
And for the temperature at the base of the H-shell
dln(T)/dln(McHe)>1
Kippenhahn & Weigert (1991)
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Caloi & Mazzitelli (1990) Sweigart (1990)
Mimicking semiconvection with Mimicking semiconvection with overshootingovershooting
Breathing pulses Breathing pulses are still found to are still found to occuroccur
Extension of mixing (by Extension of mixing (by ~0.1Hp) ~0.1Hp) in regions beyond the boundary in regions beyond the boundary of all convective regions (core of all convective regions (core and shells) forming within the and shells) forming within the He-rich core (He-rich core (~0.1Hp)~0.1Hp)
The edge of the convective core The edge of the convective core is let is let
propagate with velocitypropagate with velocity
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‹‹‹‹‹‹‹‹ ‹‹‹‹‹‹‹‹ Semiconvection Semiconvection and HRD evolutionand HRD evolution
Semiconvection Semiconvection increases central He-increases central He-burning lifetime by a burning lifetime by a factor factor ~1.5 - 2~1.5 - 2
Breathing Pulses Breathing Pulses ››››››››››››››››
(Start when Y(Start when Ycc~0.10)~0.10)
Numerical artifact ??Numerical artifact ??Parameter R2=Nagb/NhbParameter R2=Nagb/Nhb
Observations R2Observations R2~~0.140.14
Semiconv +BPs R2=0.08Semiconv +BPs R2=0.08
Semiconv noSemiconv no BPs R2=0.12 BPs R2=0.12
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Too large overshooting erases the partial mixing Too large overshooting erases the partial mixing profileprofile
SemiconvectionSemiconvection Overshooting Overshooting 1Hp1Hp
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II phase: Development of a II phase: Development of a ‘partial mixing zone’ ‘partial mixing zone’
When YWhen Yc c decreases below decreases below ~~0.7, 0.7, a ‘partial mixing’ a ‘partial mixing’ (semiconvective) zone develops (semiconvective) zone develops beyond the boundary of the beyond the boundary of the convective core.convective core.
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