Stellar evolution in a nutshell - obs

Stellar evolution in a nutshell

When perfect gaz prevails hydrostatic equilibrium implies continuous loss of energy

Star compensate for this loss either by macroscopic contraction or microscopic ones

This changes the structures and the composition of the star

Nuclear reactions in

Perfect gaz conditions

Stable

thermostate

Contraction in perfect

Gaz condition

Increases of central

temperature

Evolution of the temperature

and density at the centre

Log

Log Tc

Slope 1/3

Pgaz=PdegNR

Slope 2/3

Pgaz=PdegNR

3/2

3/51

1

e

H

k

mKT

3/5

1

eKT

m

k

H

Nuclear reactions in

Degenerate conditions

unstable

Flash/destruction

Contraction in perfect

degenerate conditions

decreases of central

temperature

Stellar evolution in a nutshell

When perfect gaz prevails hydrostatic equilibrium implies continuous loss of energy

Star compensate for this loss either by macroscopic contraction or microscopic ones

This changes the structures and the composition of the star

These processes drive the central regions in degenerate regimes

In degenerste regime: nuclear reaction unstable, contraction may lead to cooling

Hydrostatique equilibrium is for free! No long evolution

What is a massive star? 1) Ignition of carbon

2) Progenitor of core-collapse supernovae

3) Progenitor of neutron stars and stellar black holes

1) Ignition of carbon

Two important limiting masses

Mup: minimum mass for C-ignition

Mmas: minimum mass for evolving

through all nuclear burning phases

Mup

Off-centre ignition!

Mup 6.6 Msol

Maeder & Meynet 1989

A&A 210 155

Mmas

Nomoto 1984

If the One core mass

at the end of the C-burning

phase is greater than

1.37 Msol, the star

proceeds through all

nuclear burning stages

and evolves into

an iron core collapse SN

Siess 2007, A&A 476, 893

Around 8 Msol

What does happen between Mup

and Mmas (6.6 and 8 Msol)?

After the core C-burning phase, the core is degenerate.

Its mass is regulated by three competing mechanisms

1) Mass loss (decreases Mcore)

2) Activity of shell-burnings (increases Mcore)

3) Deepening of the outer convective zone (decreases Mcore)

The degenerate core never exceeds the Chandrasekhar mass

Evolution of the core into an ONe white-dwarf (SAGB)

The degenerate core mass becomes > Chandrasekhar mass

electron capture supernova type IIlikely neutron star

Convective

core

What is difference between a convective and a radiative

region??

Brunt Vaisala frequency

2) Minimum mass for the

progenitors of type II Sne (IIP)

Smartt, 2009, ARAA, 47, 63

Maximum likelihood analysis

8.5 Msol +1 -1.5Msol

3) Maximum mass for progenitors of WD?

ONe cores WD

ONe cores electron capture SNe type II NS

6.6Msol 8Msol

6Msol 7Msol 8Msol

Where is MWD max. Initial mass For progenitors of white dwarfs?

Observational determination of the maximum mass of the progenitors of white dwarfs.

Semi-empirical method

1. Detection of WD in open clusters

2. From spectroscopy determine Teff, log g

3. From the mass-radius relation, MWD

4. Age of the cluster=nuclear lifetime+cooling age

5. From the nuclear lifetime deduce Mini

Mini

MWD

1.4M

MWD

Williams, et al. 2009, ApJ 693, 355

Few stars here

MWD

~ 6.3 - 7.1Msol

One dominant systematicuncertainty comes from age

determination

See alsoKoester & Reimers 1996Weidemann 2000

In the age range of a few ten Myr.

tMS(8Msol)~40 Myr.

20% of uncertainty on the age

1 Msol uncertainty on MWD

WHERE TO IMPROVE?

Still relatively far from obtaining ages

with a precision better than 10-20%.

Meynet et al. IAU Symposium 258

What is the upper limit?

For a given hot mass star there exist a maximum value of

the luminosity called the Eddington luminosity

Msol

M

Lsol

L 41082.3

This luminosity is such that the outward acceleration given to

the matter through the interactions between photons and

electrons (through electron scattering) is equal to the gravity

at the surface of the star

An upper bound of stellar

mass

195 Msol

Lmax 7.4 106 Lsol

Milky way, NGC 3603 :

104 Msol, , age ~1.5 Myr

Masses 83 – 180 Msol

LMC, R136 :

5 104 Msol, , age ~1.7 Myr

Masses 135- 320 Msol

Crowther, Schnurr, Hirschi, Yusof, Parker, Goodwin, Abu Kassim, 2010, MNRAS, in press

STARS WITH MASSES ABOVE 150 Msol

Rotating tracks Vini/Vcrit=0.4

Mass fraction in massive stars

M > 8 Msol 14%

Stars formed between 0.1 and 120 Msol

Salpeter’s IMF

Low and intermediate mass stars

1< M/Msol < 8 25%

Very low mass stars

0.1 < M/Msol < 1 61%

IN A STELLAR GENERATION: 3/1000

with masses between 8 ans 120 Msol

PROPERTIES OF MASSIVE STARS

MASSIVE STARS AS COSMIC ENGINES

Massive stars plays a key role in many cosmic evolution processes…

rare

but very powerful emitters of

RADIATION

MASS

MOMENTUM

Short lifetimes

Mowlavi, Meynet , Maeder , Schaerer , Charbonnel , A&A 335, 573 (1998)

3 millions d’années

30 millions d’années

40

Massive stars can be seen far away in the

universe

galaxy D[kpc] mv(B superG)

LMC 46 8.5

SMC 63 9.8

M31 724 14.3

M81 3300 17.7

M100 17000 26.5

VLT limiting magnitude 28.5,

a 25 Msol star can be detected

up to distances > 70 Mpc

PHOTOMETRIC OBSERVATIONS WELL BEYOND THE LOCAl GROUP

~ 70 Mpc ~1.7% radius of observable Universe

30 Mpc

INTENSE SOURCES OF RADIATION

2/3 of the visible light of the galaxies

The Wolf-Rayet star WR224

is found in the nebula M1-67

which has a diameter of about

1000 AU

The wind is clearly very

clumpy and filamentary.

A 60 Msol

Evaporating stars

60 Msol 14 Msol

Mass fraction in massive stars

M > 8 Msol 14%

Stars formed between 0.1 and 120 Msol

Salpeter’s IMF

Low and intermediate mass stars

1< M/Msol < 8 25%

Very low mass stars

0.1 < M/Msol < 1 61%

~1% in remnants

~13% returned

~6.5% in remnants

~18.5% returned

If BH for M> 30 Msol

~ 7.4% returned

Z-dependence ?

3.5 – 4.5 % new elements

MASSIVE STARS ARE MECHANICAL STARS

(M > 40Msol at solar metallicity)

SNe! similar to ergs 10 2

years 000 500 During

1.0 ,30000

km/s 3000 ,2

1

51

2

XE

L

LLL

vvML

mechanic

mechanic

solmechanic

mechanic

Evolution of a typical massive star

Hirschi, Meynet, Maeder, 2004, A&A, 425, 649

Vrot=0 km s-1

Vrot=100 km s-1Vrot=200 km s-1Vrot=300 km s-1

B type star

B supergiant Red supergiant

After the core He-burning phase

little evolution in the HR diagram

(unless strong mass loss occurs)

Very short lifetimes for

advanced phases

H-burning

Si-burningC-burning

He-burning O-burning

Ne-burning

THE CNO CYCLE

How can we check the

models?

Studying the number of stars of different types

Measuring the surface composition of stars

New technics: asterosismology, interferometry,

neutrino emissions…

Distribution of stars: an example

Dificulty: stars must be cluster

members, sample complete, identification

All stars same age blue SG and red SG

correspond to different mass ranges

The theory

R

R

B

B

dMdM

dN

dMdM

dN

R

Bmax

min

max

min

Clusters or short starbursts

Caron et al. 2003, ApJ, 126, 1415

For field populations: mixture of ages,

ratios depend on star formation history

R

R

RSG

B

B

BSG

dMMtdM

dN

dMMtdM

dN

R

Bmax

min

max

min

)(

)(

If constant star formation rateDohm-Palmer & Skillman, 2002, ApJ, 123, 1433

CHANGE OF SURFACE ABUNDANCES

Lamers et al 2001

Przybilla et al. 2010 , A&A, 517, A38

Changes of surface abundances at the surface of MS stars

observed but not predicted by standard models

Internal mixing

Mass loss

Mass accretion from an evolved companion

Abondance anomalies present already in the cloud

from which the star formed

SURFACE ABUNDANCES

Cells of meridional circulation

GRATTON-

ÖPIK CELL

What can be the cause?

Internal mixing?

Mass loss?

Mass transfer in a close binary?

Next chapter: study of the impact of various ingredients of the models

SOME RECENT OBSERVATIONS

Surface abundances

Massive star populations

Rotational velocities

Surface magnetic fields

Mass loss by stellar winds

Enrichments at early stages

depend on Z

Depend on Z ?

Still few observations, likely important

Depends on stellar models

but also on SFH

IN A STELLAR GENERATION: 3/1000 !

Diamètre de HD 206936, 1500 Rsol

Freytag (2003)

Massey et al. 2005

Massive stars are

generous...

SN Ic

GRB MASSIVE STARS

Record redshift 6.29!

Hjorth et al. 2003, Nature, 423, 847

Stanek et al. 2003, ApJ, 591, L17

Ekin=4 x 1052 ergs

0.35 Msol of 56Ni

Ejecta ~8 Msol

Mass 25-30 Msol

SOURCES DE PHOTONS IONISANTS

REIONIZATION DE L’UNIVERS A HAUT REDSHIFT

Hester et al. 2004

z 17.6 z 15.5 z 13.7

Norm

al sta

rs

Mc

~ 1

Ms

un

Massiv

e s

tars

Mc

>>

1 M

su

n

Ciardi, Ferrara & White 2003

REIONIZATION OF THE UNIVERSE AT HIGH REDSHIFT

OBSERVATIONAL EVIDENCES FOR MIXING

• Extended cluster MS Maeder, 76; Mermilliod et al. 93

• ON stars Walborn, 76, 2002;

Heap & Lanz 2003

• Fast rotators with He, N excesses Lyubimkov 91-98; Daflon et al. 99, 01

Herrero et al. 92; Villamariz et al. 02

• He, N excesses in B, A, F supergiants Gies & Lambert 92;

Lennon 92, 2002

Venn 95, 2002

Venn and Przybilla 2003

• Stronger He, N excesses in SMC supergiants Venn 95, 2002

• He, N excesses in SN 1987A Fransson et al. 89

• Boron depletion in rotating B-stars Fliegner et al. 96; Venn et al. 96, 2002

• Transition WN/WC stars Langer 91; Crowther 95, 02;

Morris et al. 99

• Blue/ Red supergiant ratios at various Z Langer & Maeder 96;

Maeder & Meynet 2002

THE CASE OF BORON IN MASSIVE STARS

AN INTERESTING CLUES FOR A SMOOTH INTERNAL MIXING

Mendel et al. 2006, ApJ, 640, 1039

Venn et al. 2002, ApJ, 565, 571

Boron destroyed by p-captures for T < 6 106 K, T reached only

about 1 Msol down in hot stars

Shallow mixing sufficient to destroy it,

but would be insufficient to change the abundances of other

elements as for instance Nitrogen

Spectroscopy of the B III feature at 2066 Angstroem

IUE archive, HST (STIS, GHRS)

B depletion without N enrichment would be sign of a shallow

mixing

A signature of a smooth internal mixing is the following

B

N

B depletion before N enhancements

B depletion without N enhancements

In hot stars

Boron lifetime

< 104 years

Mendel et al. 2006

THE SAMPLE

B stars in

Young clusters

of the solar

neighborhood

Masses between

8 and 14 Msol

Mass loss

very weak

The results

Significant Boron depletion

BEFORE

Significant N

enrichment

Morel et al 2006, A&A, 457, 651

Enhancement factor of the N/C ratio with respect to the Sun (Asplund et al. 2005)

2.6

1.5

3.5

1.73.7

2.5

Meynet 2002

Crowther et al. 2006

Galactic Early B Supergiants

N/C 10 X solar value

N/O 5 X solar value

Greater N enrichment for

Lower Teff

Venn & Przybilla 2003

Max/ini N/H =40

Max/ini N/H =8

Log (N/H)+128.88.48.07.67.26.86.4

Number

of stars

Mokiem et al. 2006, A&A, 456, 1131

31 O- and B-type stars

in SMC (21 in NGC 346)

with VLT-FLAMES

MIXING IN MASSIVE STARS

Observational evidences

Solar neighborhood: typical values N/H in excess

with respect to initial value by about a factor 2 at the

end of the MS phase

Values in excess by a factor 8-10 are possible

SMC: for stars above 20 Msol values in excess by a factor

30 seem to be relatively easily reached

At low metallicities, mixing appear to be more efficient

Internal mixing

Mass loss




SURFACE ABUNDANCES


GRATTON-

ÖPIK CELL

Hamann et al., 2006, A&A, 457, 1015

THE WOLF-RAYET STARS

Gemini obs.

of Wolf-Rayet

stars in the

starburst galaxy

IC10

Crowther et

al. 2003

AA, 404, 483

Crowther

et al. 2003

WR = bare cores of initially massive stars (M > ~40 Msol) whose

original H-envelope has been removed by stellar winds or

through Roche lobe overflow

N/Cini=0.25

N/Oini=0.11

N/Ccno=50

N/Ocno=10

X 200

X 100

CNO in massive stars: N produced at the expense of O and C

He burning: C and O is created. N destroyed.

Observations by Crowther P.A., Smith L.J., Willis A.J. 1995, A\&A 304, 269


Neon too high!

With the new solar abundances.

Internal mixing

Mass loss




SURFACE ABUNDANCES


GRATTON-

ÖPIK CELL

Galaxy Z WR/O WRRLOF/WR

SMC 0.002 0.021 0.98 ± 0.32

LMC 0.006 0.05 0.41 ± 0.13

Milky Way 0.02 0.104 0.20 ± 0.06

Foellmi et al. 2003ab

Obs. Theory Obs. Binaries

0.40

0.30

0.40

Tuthill et al. (1999)

Tuthill et al 2006Pinwheels in the Quintuplet Cluster

A NEW WINDOW ON THE LATE STAGES OF MASSIVE

STAR EVOLUTION

Low Mass X-Ray Binaries

Part of matter ejected at the time of the formation of the

compact object pollutes the companion

Nova Sco 94 (Israelian et al. 1999; Brown et al. 2000; Podsiadlowski et al. 2002)

A0620-00 (González Hernández et al. 2004).

Centaurus X-4 (González Hernández et al. 2005).

G and K metal-rich dwarf stars

Bodaghee et al. (2003)

Feltzing & Gustafsson (1998)

Internal mixing

Mass loss




SURFACE ABUNDANCES


GRATTON-

ÖPIK CELL

Extremely metal poor C-rich stars

Frebel et al. 2005 [Fe/H]=-5.4

Plez and Cohen 2005 [Fe/H]=-4.0

Christlieb et al. 2004; Norris et al. 2001

Depagne et al. 2002; Aoki et al. 2004

HE0107-5240

0.8 Msol

Z = 0.0000001 Z = 0.001

Z = 0.020

Initial composition

First Dredge-up

[Fe/H]=-5.3

[C/Fe]=+4.0

[N/Fe]=+0.0

[O/Fe]=+2.3

Solar ratios for the other elements

After first dredge up

[Fe/H]=-5.3 Christlieb et al. 2002

[C/Fe]=+4.0

[N/Fe]=+2.3

[O/Fe]=+2.3 Bessel et al. 2004

[Fe/H]=-5.3

[C/Fe]=+4.0

[N/Fe]=+2.7

[O/Fe]=+2.3

[N/H]=+2.3

[N/H]=+1.0

[N/H]=+2.8

[N/H]=+0.0

Nitrogen CAN BE PRODUCED IN SITU

BUT Carbon AND Oxygen

HAVE TO BE EXPLAINED

Umeda and Nomoto 2003; Limongi et al. 2003

Meynet et al. 2006; Hirschi 2007

SURFACE MAGNETIC FIELDS

Donati et al. 2006

t Sco

MAGNETIC FIELDS IN MASSIVE STARS

A few dozen He-peculiar stars

Only 7 OB stars have been found to be magnetic

Ref Sp. T. Vsini

Km/s

Prot

days

M

Msol

Incl.

Deg.

b

Deg.

Bpol

G

HD191612 (6) 538 45 ~1500

Q Ori C (1) O4-6V 20 15.4 45 45 42+-6 1100+-100

bCep (2) B1IVe 27 12.00 12 60+-10 85+-10 360+-40

t Sco (7) B0.2V 41 ~500

V2052 Oph (3) B1V 63 3.64 10 71+-10 35+-17 250+-190

zCas (4) B2IV 17 5.37 9 18+-4 80+-4 340+-90

wOri (5) B2IVe 172 1.29 8 42+-7 50+-25 530+-200

He-peculiar B1-B8p 0.9-22 <10 1000-

10000

(1) Donati et al. 2003 (2) Henrichs et al. 2000 (3,4,5) Neiner et al. 2003abc, (6,7) Donati et al. 2006ab

b Angle between the magnetic axis and the rotation axis

Only 2 magnetic

O star known

Question: are these values compatible with magnetic fields

observed in pulsars?

2

2

)/(/

.

rrBB

constBr (10 km/5 Rsol)2 x 1012 G ~ 10 G.

1012 GPulsars

Answer: observed magnetic are one-two orders of magnitude

higherMore compatible with progenitors of magnetars 1015 G

Question: may the observed values have an impact on the wind?

2/

8/)

2

2

v

Br

p

ud-Doula & Owocki (2002)

Answer: YES. For early-type stars, >1 for B~ 50-100 G

if > 1 wind behavior

Ref Sp. T. He I C II N II O II

bCep (2) B1IVe 0.09 (1.2)

+-0.06

V2052 Oph (3) B1V 0.32 (2.1)

+-0.05

-0.13 (0.7)

+-0.04

0.10 (1.3)

+-0.06

-0.31 (0.5)

+-0.11

zCas (4) B2IV 0.11 (1.3)

+-0.06

-0.05 (0.9)

+-0.09

0.41 (2.6)

+-0.10

-0.09 (0.8)

+-0.14

wOri (5) B2IVe 0.00 (1.0)

+-0.01

0.00 (1.0)

+-0.07

0.26 (1.8)

+-0.10

-0.09 (0.8)

+-0.06

All magnetic B stars appeared to have some abundance anomaly

Log [number nuclei N in star/number nuclei N in the Sun]Grevesse & Sauval 1998

N/C=1.9

N/C=2.9

N/C=1.8

STAR POPULATIONS

RSG BSG

When the metallicity decreases, models predict that the number ratio of blue

to red supergiants increases.

DO WE OBSERVE THIS TREND ?

Eggenberger et al. AA, 386, 576 (2002); Langer and Maeder AA 373, 555 (1995)

NUMBER RATIOS OF

MASSIVE STARS

IN NEARBY GALAXIES

M31 0.035 0.24 0.44 1.7

6-7.5 0.029 0.21 0.55 --

7.5-9 0.020 0.104 0.48 ~1

9.5-

11

0.013 0.033 0.33 --

M33 0.013 0.06 0.52 ~4

LMC 0.006 0.04 0.20 --

6822 0.005 0.02 -- 8.3

SMC 0.002 0.017 0.11 --

1613 0.002 0.02

GALAXY Z WR/O WC/WR RSG/WR

Conti & Maeder’94;

Massey ‘02

OBSERVATIONS OF WR POPULATION IN CONSTANT SF REGIONS

LES GALAXIES WOLF-RAYET

Fernandes, de Carvalho, Contini, Gal, MNRAS 355, 728 (2004)

NO NWR

5660

+/- 600

58500

+/-17000

1800

+/- 390

2000

+/- 150

NWR

700

+/- 360

40000

+/- 16000

660

+/- 330

450

+/- 220

AGES, ENVIRONNEMENTS, SNe

Fernandes, de Carvlho, Contini, Gal, MNRAS 355, 728 (2004)

NO NWR

5660

+/- 600

58500

+/-17000

1800

+/- 390

2000

+/- 150

NWR

700

+/- 360

40000

+/- 16000

660

+/- 330

450

+/- 220

IIIZw107

MRK 475

MRK 1271

NGC4385

Meynet, A&A 298, 767 (1995)

WC/WN

Massey and Johnson 1998

Schild et al. 2003


Hadfield et al. 2005

Classification of Supernovae

Type Ia Ib Ic II

Spectrum No Hydrogen

Silicon

No Hydrogen

No Silicon

Helium

No Hydrogen

No Silicon

No Helium

Hydrogen

Physical

Mechanism

Nuclear

explosion of

low mass star

Core collapse of evolved massive star ( may have lost its hydrogen or even helium

during red-giant evolution)

Light

Curve

Reproducible Large Variations

Neutrinos Insignificant ~100 times the visible energy

Compact

remnant

None Neutron Star (typically appears as a pulsar)

Sometimes Quark star?

Sometimes Black hole ?

Observed Total : ~ 2000 as of today (nowadays ~ 200 / year )

Prantzos and Boissier (2003)

Stellar Population Model

Theory: A. Renzini (1981)

The Fuel Consumption Theorem

The theoretical spectra of Stellar Populations: composing spectra of

individual stars weighted by the fuel

Individual Stars

Integrated over all starsMaraston 1998

Abundances in the winds of Wolf-Rayet stars

WN XN~0.01 XC~0.0005

WN/WC XC~0.05XN~0.01

WC/WO XC~0.20-0.55 XO~0.05-0.10 XNe~0.01

See references in review of Crowther ARAA 45 (2007)

4/p

Huang and Gies 2006

461 OB stars in 19 young clusters

Average correction for Vsin i 4/p = 1.27

Masses between 3 and 15 Msol

much more important than in low mass stars ....

ORIGIN OF LONG SOFT GRB ?

~3 per day in the observable Universe

~5 core collapse SNe per second

Average factor of beaming of 300

GRB rate ~0.2% SN rate

EARLY UNIVERSE

Reionization Abondances at the surface of extremely

metal poor stars

z 17.6 z 15.5 z 13.7

No

al

sta

s

Mc

1 M

su

n

Ma

ss

ive

sta

s

Mc

>>

1 M

su

n

Ciardi, Ferrara & White 2003

At high redshift In the local Universe

Injection of mechanical energy

Winds 1051 ergsSN 1051 ergs

Blue

Supergiant

Sher 25

with a ring.

Sher 25

87A

Ratio Solar Sher 25

N/C 0.25 26.3

N/O 0.12 0.36

Smartt et al. 2002

Incompatible with the star having a

Previous RSG phase

« … the radiation observed to be emitted must workits way through the star, and if there were too muchobstruction it would blow up the star. »

Eddington 1926

v

vv

A differentially expanding radiation-driven

spherical shell in the wind

rrmmt

momentumabsorbedgrad

p 24,

i

L

L

c

L i

nn

n n

width andfrequency of

i line strongby abs.Fraction

star by theprovided

mom. Tot.

i

t

mom. abs.

The velocity gradient due to the

differential expansion is so large that

the line bandwidth across the shell is

determined by the Doppler formula

(Sobolev approximation)

iic

nn

n

i

L

L

c

L i

nn

n n

width andfrequency of

i line strongby abs.Fraction

star by theprovided

mom. Tot.

i

t

mom. abs.

iic

nn

n

dr

dv

rL

L

c

Lg

effN

i

irad

p

nn

22 4

1

rrmmt

momentumabsorbedgrad

p 24,

dr

dv

rL

L

c

Lg

effN

i

irad

p

nn

22 4

1

Radiative acceleration is proportional to the luminosity and to the velocity gradient

effNResults from the sum over all lines and is interpreted as the number

of effectively acting strong lines.

Equation of motion of the stellar wind in stationary regime in the supersonic

region

Thomson

2

y)(stationar 0

)1(

gg

radr

GMg

dr

dvv

t

v

ep

142

2

effNc

LMvrM

effLN

c

drdv

GM 2

/

4)1( pe

~0.1 for luminous

Galactic OB stars

Fukuda, PASP, 94, 271, (1982)

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