Today in Astronomy 142

Star clusters and stellar evolution: The final stages of stellar evolution

Observations of stellar evolution: the Hertzsprung-Russell diagrams of open and globular stellar clusters

21 February 2013 Astronomy 142, Spring 2013 1

Open clusters M 35 (left) and NGC 2158 (right) in Gemini (Dieter Willasch).

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Late stages of stellar evolution

After the maximum mass fraction of the isothermal helium core is exceeded, the star enters the… Red giant phase (moving up in H-R diagram) Core collapses: central density and temperature rise. Convection zone extends inward (dredge-up). Stellar radius increases dramatically, from sharp increase

in radiation pressure from interior. Radiation pressure now dominates support against the star’s weight.

Core temperature reaches 108 K, and the triple-α process, begins helium fusion. The onset of this process is very rapid in stars with leading to a phenomenon called the helium flash.

2 ,M M≥

4 8 0 4 12 02 4 0 2 6 03 He Be* He C 2 ,γ γ→ + + → +

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Late stages of stellar evolution (continued)

The horizontal branch is the phase after triple-α onset. Core helium burning, shell hydrogen burning. The core is

on the helium main sequence. (You will explore the properties of the helium main sequence in Homework #4.)

0.1

1

10

100

1 103

1 104

Globular cluster M3Eclipsing binaries

Log of effective temperature (K)

Lum

inos

ity (s

olar

lum

inos

ities

)

Horizontal branch

Red giants

Subgiants

4.5 4 3.5 (Hydrogen) main sequence

Nomenclature: spectral type

The “Harvard” spectral type of a star is a classification based upon the strength of absorption lines of several molecules, atoms and ions in its spectrum (Annie J. Cannon, 1901). Generally the types A-M correspond to steady decreases

in the strength of hydrogen lines; type O “weaker” still (emission lines instead of absorption), P for planetary nebulae, Q for miscellany. N was not used.

Years later Cecilia Payne (1925) showed that the spectral-type sequence OBAFGKM is a sequence of decreasing temperature, roughly 35000 – 3500 K. More recently types L, T and Y have been added for brown dwarfs.

Examples: Vega is an A0 star, the Sun a G2 star, Pollux a K2 star, and Betelgeuse an M2 star. 21 February 2013 Astronomy 142, Spring 2013 4

Nomenclature: luminosity class

At Yerkes Observatory in the 1940s, Morgan, Keenan and Kellman (1943; MKK) added another dimension to classification by introduction of luminosity classes. V: main sequence, or dwarf, stars. IV: subgiants, brighter than V by a magnitude or two for

the same spectral type. III: normal giants, another few magnitudes brighter II,I: bright giants and supergiants, yet another few

magnitudes brighter. VI, VII: subdwarfs and white dwarfs.

Examples: Vega is an A0V star, the Sun is G2V, Pollux is K2III, and Betelgeuse is M2I.

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Spectral type, Te, and luminosity class

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0

5000

10000

15000

20000

25000

30000

35000

40000

45000

O5 B0 B5 A0 A5 F0 F5 G0 G5 K0 K5M0 M5

Effe

ctiv

e te

mpe

ratu

re (K

)

Spectral type

VIIII

-15

-10

-5

0

5

10

15O5 B0 B5 A0 A5 F0 F5 G0 G5 K0 K5M0 M5

Abs

olut

e bo

lom

etric

mag

nitu

de

Spectral type

IIIIV

Data from Schmidt-Kaler 1982

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After the horizontal branch

Low mass stars (those with ): Slowly an “isothermal carbon-oxygen core” forms in the

center as the helium fuel is exhausted. In these stars, however, there is not enough weight to

overpower degeneracy pressure, so the core can’t collapse and reheat to ignite carbon-oxygen fusion.

Results: • H/He burning of outer layers of core, ejection of most

of the rest of the star, and formation of a planetary nebula. This phase lasts a few thousand years, during which the ejected gas drifts away.

• The core becomes a carbon-oxygen white dwarf with mass and initial temperature ~108 K. Lasts forever unless it has a close stellar companion.

2M M<

0.6M M≈

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After the horizontal branch (continued)

Massive stars (those with ): Asymptotic giant branch (AGB, or supergiant) evolution. Repeated core collapse - fusion reignition - nuclear fuel

exhaustion occurs, including silicon fusion to produce iron-peak elements.

Radius and luminosity steadily increase, Te decreases. Each of the successive fuel exhaustions is faster than the

last. For a star, • hydrogen burning (main sequence) lasts 107 years • helium burning (horizontal branch) lasts 106 years • carbon burning lasts 300 years • oxygen burning lasts 200 days • silicon burning lasts 2 days !

2M M>

20M

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Simulation of stellar evolution from main sequence to AGB

Terry Herter’s (Cornell) AST 101/103 site, the source of the nice binary-star simulations we’ve used, also has a stellar-evolution simulator. Have a look here.

Results from Terry’s simulator, for stars with The tracks take 5.6 Gyr for the lighter star and only 14 Myr for the heavier one.

1.5 and 15 .M M=

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What happens when all the nuclear fuel is gone?

Most stars: During the burning of heavier elements, and radiative

support of the stellar envelope, stars tend to be hydrodynamically unstable, leading to the loss of large fractions of stars’ masses: • Oscillations: note that evolution takes stars across the

instability strip (lecture, 7 February), which is nearly vertical at effective temperature ~ 5,000K.

• Stellar winds (see this week’s Recitation, and HW#4). This can keep a star’s core mass below the SAC limit, and

the final states of the star are just like that of less massive ones: planetary nebula phase and white dwarf remnant.

2M M>

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What happens when all the nuclear fuel is gone? (continued)

The most massive stars Mass loss insufficient to keep core in white dwarf range:

further collapse and neutronization. • Maximum mass of iron white dwarf

When the collapsing core reaches tens-of-km dimensions, neutron degeneracy pressure sets in, and this can stop or slow the collapse.

However, since the collapse has been from white-dwarf dimensions to neutron-star dimensions, infalling material from the star’s envelope is going very fast: a large fraction of c. It bounces off the stiffened neutron-degenerate material and blows up the rest of the star.

( )8 :M M

1.26 .M=

Core collapse and Type II supernovae

This event is called a core collapse, or type II, supernova. • How did we get to type

II before type I? We’ll get to that later in the course, when we talk about the extragalactic distance scale.

Remnant: a neutron star or, more rarely, a black hole, depending upon core mass.

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Artist’s conception of a SNII and the formation of a supernova remnant (Chandra/CfA/NASA). ( 1.3 2.2 )M M= −

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A SNII forms from a dead, massive star (not drawn to scale)

Star: 6 M

, 107 km circumference Core: 1.4 M

, 105 km circumference

Core: 104 km circum-ference. Electrons and protons begin combining to form neutrons.

2 years

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A SNII forms from a dead, massive star (continued)

Core: 75 km circumference, neutron degeneracy pressure sets in.

Core: 104 km circumference. Electrons and protons begin combining to form neutrons.

1.2 seconds

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A SNII forms from a dead, massive star (continued)

Core: 75 km circumference, neutron degeneracy pressure sets in. This makes the core very stiff.

Outside of core: still collapsing, moving inwards at about 1010 cm/s. Bounces off stiff core.

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A SNII forms from a dead, massive star (continued)

Core: Still 75 km circumference, it is now stable.

Outside of core: the rebounding outer-star material explodes the rest of the star. Energy comes from bounce, and from gravitational energy of core.

A few seconds

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A SNII forms from a dead, massive star (continued)

Neutron star

About a day

Expanding supernova shell. Very, very bright for about a month after explosion (can outshine rest of galaxy!). We’ll talk about supernova remnants two lectures hence.

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Supernova 1987A in the Large Magellanic Cloud

Before (top; follow the arrow) and after (bottom; guess where) the explosion.

SN1987A was the first supernova for which we knew the progenitor star, and was the most recent SN that could be seen with the naked eye. It was a SNII.

David Malin, Anglo-Australian Observatory

Supernova light curves

Type I and II supernovae are distinguished observationally by the shape of their luminosity as a function of time (light curve). SNe I decline in

brightness by about 0.1 mag/day and reach peak absolute magnitudes in the V = -18 — -19 range.

SNe II decline more slowly ( 0.05 mag/ day) and have smaller peak absolute magnitudes (V ~ -17).

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Wikimedia Commons

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Observation of stellar evolution: star clusters

Stars tend to form in clusters, with all members within 1-2 My of having the same age. Open clusters (young): low density, irregular, lots of blue

stars, low random velocities (few km/sec), hundreds to thousands of stars, not always gravitationally bound. Archetypes: Pleiades (M45), Hyades, h and χ Persei.

Globular clusters (old): high density, spherically symmetrical, few blue stars, higher random velocities (tens of km/sec), millions of stars, gravitationally bound. Archetypes: ω Centauri, M3, M13, 47 Tucanae.

Star clusters are very useful for studying stellar evolution and for determination of distance scales in the universe.

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The (observer’s) H-R diagram for clusters

The plot of apparent magnitude in the V band (backwards) against the color index B-V is the classical Hertzsprung-Russell diagram. Plotted in this way such diagrams will resemble our previous logarithmic plots of luminosity vs. effective temperature (backwards).

V

B-V

Main sequence L

Te

Main sequence

Evolution of stars in a cluster

Because the measurement of an HR diagram is a snapshot in stellar age terms – made in an instant – the HR diagram of a cluster does not resemble the tracks in L-Te space followed by individual stars as they age. Instead there will tend to be a turnoff at some point on

the main sequence, above which there are no main sequence stars left. All the stars originally on the upper main sequence have moved elsewhere in the HR diagram, with identifiable collections only in the longer-lived stages like the RGB and the horizontal branch.

Nice animation of cluster evolution by Leslie Tomley (San Jose State): look here. Does white-dwarf cooling way too fast but is otherwise accurate.

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Evolution of stars in a cluster (continued)

Bolo

met

ric

mag

nitu

de

Bolo

met

ric

mag

nitu

de t = 0

t = 107 yrs

t = 108 yrs

t = 109 yrs

t = 1010 yrs

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0 .5 0 0 .5 1 1 .5 2

0

5

10

15

20

Open clusters: the Pleiades (M45)

V

B-V Image: David Malin, Anglo-Australian Observatory

HR diagram of Pleiades X-ray sources (Stauffer et al. 1994)

Only main sequence

The Pleiades lie about 135 pc away and are about 110 Myr old.

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0

2

4

6

8

10-1 0 1 2

B-VM

V

Open clusters: the Hyades

Image: Hermann Gumpp; HR diagram: Hipparcos (ESA)

A few giants

MS

Hyades Pleiades

The Hyades are the closest open cluster (43 pc) and are about 660 Myr old.

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Open clusters: M67

8

10

12

14

16

18

20

220 0.5 1 1.5 2

B-V

V

RG

SG

MS

Image: Sharp and Hanna (NOAO/AURA/NSF); data: Montgomery, Marschall and Janes (1993)

M67 is about 900 pc away and is about 4×109 years old; it’s the oldest known open cluster .

MS turnoff

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Globular clusters: M3

1 0.5 0 0.5 1 1.5 2

12

24

V

B-V

MS

SG

RG

HB AGB

Blue stragglers

MS turnoff

Image: Peter Challis (Harvard-Smithsonian CfA); data: Ferarro et al. (1997).

Like all Galactic globular clusters, M3 is about 12000 Myr old. It lies about 10400 pc away. RR Lyrae stars not plotted (note gap in HB).