Download Fifth - Department of Physics and Astronomy

ASTR377:
A six week marathon
through the firmament
Week 5, May 17-20, 2009
by
Orsola De Marco
[email protected]
Office: E7A 316
Phone: 9850 4241
Overview of the course
1.
Where and what are the stars. How we perceive them, how we
measure them.
2.
(Almost) 8 things about stars: stellar structure equations.
3.
The stellar furnace and stellar change.
5.
Stellar death: stellar remnants; PN, WDs, NS, BH!
6.
When it takes two to tango: binaries and binary interactions.
What death becomes her?
• Depending on the main sequence mass, a star
will end in a different way.
• M<0.08Mo no H burning – BD [L or T dwarfs].
• 0.08Mo < M < 0.5Mo no He burning. Stars
become He WDs.
• 0.5Mo < M < 5Mo no C burning: stars become CO
WDs.
• 5Mo < M < 7Mo yes C burning to Ne and Mg:
stars become ONeMg WDs.
• M > 7Mo burn all the way to Fe: these stars go
through a type II SN and become NSs or BHs.
What happens after the AGB?
• Spectra of planetary nebulae indicate that the
“shell” is expanding with speeds of 20-30 km/s.
The central star is often visible inside the PN.
Jacoby, De Marco & Sawyer 1998
What happens after the AGB?
• Knowing angular size, distance and
speed of expansion and assuming that
the speed did not change since the time
of the material ejection, we can
determine
• What is wrong with the assumption
above (see in a few slides)?
Angular size and distance
3476 km
?
D
Base
Hypotenuse
= sin 
Moon diameter
= sin 
D
Bob O’Dell American (alive)
The PN clock
PN “ignition”: log Teff ~ 4.4
• The central stars of
PN on the HRD: they
sit in a locus and
have increasing ages
along the red arrow.
• Conclusion: the star
evolves to the blue in
a very short time.
AGB
logTeff=3.5
Youngest
PNe
Oldest PNe
O’Dell 1968
At the end of the AGB
•
•
•
•
•
Core mass growing because of the He
shell burning. This generates an increase
in L.
Two shell sources, H (out) and He (in).
Every time a star has multiple shell
sources, all burning outwards through the
fresh fuel supply, there is instability.
The relative speed at which they burn out
can create an instability where the helium
shell L increases with no release of
pressure, leading to even more L.
Eventually the L pushes the envelope
out, the H shell extinguished and the He
burning rate decreases, L gets out, then
the entire star returns to equilibrium.
At the end of each thermal shell flash the
envelope convection zone extends
downwards and dredges up the results of
He burning: C and O. This is the third
and last dredge up chance a star has and
makes of AGB star the C factories of the
Universe.
http://outreach.atnf.csiro.au/education/senior/astrophysics/stellarevolution_postmain.html
Departing the AGB
• When the AGB H envelope
mass drops below a few
times 10-4 Mo the star loses
its equilibrium and contracts.
• The contraction and the
increased transparency of
the envelope result in a fast
increase of the effective
temperature.
• All the while the H or the He
shell source are still burning.
• Eventually the burning stops
and the star cools on the WD
cooling track.
Vassiliadis & Wood 1994
The PN phase
• As the star heats a fast
but this time tenuous
wind sweeps AGB
wind material up and
creates a shells.
• When the photosphere
of the heating star
passes the ~25,000-K
mark the swept up
shell is ionized and can
be seen in forbidden
lines.
Animation from the Space Telescope Science Institute
What do Planetary Nebulae
look like in the sky
3 arcmin
Planetary Nebula shapes:
round….
Abell 39; WIYN image; G. Jacoby
Planetary Nebula shapes:
“elliptical”….
The Helix nebula; Spitzer image; K. Su
Planetary Nebula shapes:
“elliptical++”….
The Cat’s Eye nebula; HST image
OH321.8+4.2; Bujarrabal; HST
PN shapes and shaping
• Young PN and pre-PN (shining from shocks not
from radiative ionization) are always non-round.
2.5 pc
NGC6543 HST/NOT [OIII]/[NII]/Ha. (P. Harrington, R. Corradi)
Planetary nebulae as we teach them
How do PN form?
CSPN
post-AGB
AGB
3.4
WD
Empirically shown to happen, theoretically unexplained
Planetary nebulae as we teach them
How do PN form?
Envelope mass < 10-3 to -4 Mo: departure from the AGB.
Fast thin post-AGB wind compresses the super-wind.
CSPN
post-AGB
AGB
3.4
WD
Kwok 1982; Balick 1987
Planetary nebulae as we teach them
How do PN form?
Post-AGB star heats up T>25,000K: ionized PN
A39: a well behaved PN
CSPN
post-AGB
AGB
WD
… but how do PN acquire their shapes?
… but how do PN acquire their shapes?
Relatively-fast rotation during AGB super-wind …
Garcia-Segura et al. 2003
… but how do PN acquire their shapes?
… and/or magnetic fields
Garcia-Segura et al. 2003
… but how do PN acquire their shapes?
… result in circumstellar material
with an equatorial enhancement
… but how do PN acquire their shapes?
When the star heats up, on its way to becoming a white dwarf, a
fast wind rums into the previously-ejected gas.
… but how do PN acquire their shapes?
Let’s zoom out …
… but how do PN acquire their shapes?
Let’s zoom out …
… the lobes perpendicular to the plane of the disk continue to expand.
… but how do PN acquire their shapes?
When the star heats up
the gas “shines” and we see the bipolar PN
… but how do PN acquire their shapes?
The problem: giant stars do not rotate fast enough.
How can a companion spin up a giant?
Animation from the Space Telescope Science Institute
White dwarfs: key properties
• Small – about
Earth radius.
• Not (all) white:
some are very
cool (~4000K).
• Super dense.
Srius A and B; HST/FGS image
Friedrich Wilhelm Bessel German 1784-1846
History of WDs
• 1844 Friedrich Bessel
discovers that Sirius
changes position and the
change is not due to
annual parallax…. There
must be an unseen
companion.
• In 1864 the companion
was found by Alvan
Graham Clark. It was
25,000K bit only 10,000th
the luminosity of the Sun.
What does this mean?
• A better solution of the
orbit gave us Sirius B’s
mass (0.9Mo) This meant
that this star was much
denser.
Ralph H. Fowler, UK 1889-1944
History of WDs
• A better solution of the orbit
gave us Sirius B’s mass
(0.9Mo) This meant that this
star was much denser.
• In 1917 Adriaan van den
Maaren discovered another
WD (a single one).
• It was not till quantum
mechanics that Ralph Fowler
determined that degeneracy
pressure was supporting the
WDs.
• Eventually Chandrasekhar
determined the mass limit
above which WDs cannot
exist (among many many
other things that got him the
Nobel Prize)
WD density and pressure
• You can determine the WD density from
a value of its mass and radius (which for
Sirius B come from observations).
• You can then determine a value for the
central pressure of the WD using the
equation of hydrostatic equilibrium.
The mass-radius relation for WDs
aka: the more you have the smaller it is!
1/ 3
 M 
R  0.01 R

0.7M 

(No demonstration)
The Chandrasekhar limit for WDs
• As you increase the mass of a WD the radius decreases,
the density increases and eventually the electrons speeds
approach the speed of light, i.e., they become relativistic.
• By equating the core pressure for a star in hydrostatic
equilibrium, to the pressure for relativistic electrons we
see that there is no radius for which the star will be in
equilibrium: the star just collapses.
• From the same equation we can derive the mass of such
star. For masses larger than the Chandrasekhar mass
limit, the star collapses.
M  1.4 M
for stability
(Board demonstration)
Neutron Stars: formation
• Si-burning adds Fe to the core which increases in mass.
And contracts under its own weight.
• Electron degeneracy provides the pressure but when the
electrons become relativistic at the Chandrasekhar limit
the star collapses in a free-fall time (<1 second).
• Protons and electrons combine to form neutrons and
neutrinos (the neutrinos take energy out).
• Collapse is halted by neutron degeneracy pressure.
• We have a neutron star.
• A neutron star is not technically a star because it is not
gaseous and it is not powered by thermonuclear
fusion….
Neutron stars: radius
• The mass radius relations for WD is the
same as that for neutron star except that
the neutron mass is in place of the
electron mass.
RNS
RWD
me M WD 



mn  M NS 
1/ 3
Neutron stars: escape velocity
• The escape velocity from a NS is about
½ of the speed of light.
2/3
2GM 1/ 2


M
8
-1
v esc  
 R 
  2 10 m s 1.4 M 

 



Neutron stars: luminosity
• When first formed a NS has T~106K. So,
despite their small radius the luminosity is not
terribly low.
• Using Wien’s law we see this spectrum peaks at
~30A or 400 eV in the X-ray range.
• But there is another type of radiation from NS.
RNS  TNS 
     L   0.2 L 
 R   T 
2
LNS
4
Neutron stars: pulsars
• As the star collapses conservation of angular momentum makes
it spin at about 0.1c, implying a rotation period of the order of
milliseconds.
• Magnetic flux is also conserved such that the surface B fields is
intensified.
• The rotating B field creates an E field that rips charged particles
from the surface of the star, which later get beamed by the B
field and ejected at the poles.
• They were discovered during
a radio survey of the Galaxy and
the first one was named LGM-1.
• There are about 500 known
pulsars and considering the
selection effects, there must
be a lot more.
Neutron stars:
pulsars
• Why are pulsars not WDs?
• If the periodicity was
due to pulsations, the
material pulsating
would have to be very
dense indeed, much
denser than a WD can be.
• If it were a rotating WD, the rotation
speed at the periphery of the WD would
have to be 100c.
• Pulsars are often seen in the middle of
SN remnants.
John Wheeler American 1911-2008
BH History
• Postulated by John
Mitchell 200 years ago as
objects with such high g
that not even light could
escape (he called them
black stars).
• Calculated by
Schwarzschild in 1916, as
soon as GR was invented.
• Died shortly after at the
front of some disease–
very sad.
• John Wheeler called them
Black Holes
Karl Schwarzschild German 1873-1916
Black Holes:
the Oppenheimer-Volkov limit
• As for the Chandrasekhar limit there is a
. It is about 3Mo – hard to
calculate because of the effects of the strong
nuclear force.
Black Holes:
the Schwarzshild radius
• By equating the escape velocity to the
speed of light you see the size of a
BH’s event horizon.
• Nothing going inside can emerge, not
particles, not light, not information.
RSch
 M 
2GM 
  2  3 km  
 c 
M 
BH: tidal ripping
• At what distance can a
BH rip a person
approaching feet first?
(assuming that a force
that can rip a person is
about 16 tonnes)
• For a 2000 Mo BH you
will be ripped at exactly
the Schwarzshild
radius.
GMm 
F   3 l
 R 
1
3
 M 
Rrip  435   km
M 
compare with
 M 
RSch  3 km  
M 
2
3
 M 
Rrip
 160  
RSch
M 
Observations of BHs
• V404 Cyg – binary
system where a K0
star donates mass to
a dark object via an
accretion disk.
• The BH is revealed
by the orbital motion
of the companion or
the luminosity of the
accretion disk.