Download A Supermassive Black Hole in the Andromeda Galaxy

A Supermassive Black Hole
in the Andromeda Galaxy
1. Structure of the M31 center
2. Dynamics of the blue nucleus P3
3. HST ACS imaging of P3
4. Is the Dark Object a Black Hole?
Ralf Bender
Max-Planck-Insitute for Extraterrestrial Physics Garching
and Observatory of the Ludwig-Maximilians-University Munich
in collaboration with John Kormendy, Tod Lauer, Jens Thomas et al.
Astrophys. J. 631, 280, 2005 and Astrophys. J., in prep., 2008
Apparent Size of Black Hole Event Horizons:
• typical Galactic stellar mass Black Hole: 0.001 µ arcsec
Cyg X-1
Apparent Size of Black Hole Event Horizons:
• typical Galactic stellar mass Black Hole: 0.001 µ arcsec
• Galactic center Black Hole:
8
µ arcsec
Star S2:
Vmax ~ 5000 km/s,
Dmin ~ 18 Billion km
Black hole:
MBH ~ 4 Million Msun
RS
~ 9 Million km
Genzel et al. 1992-2004, Ghez et al. 1996-2004
Apparent Size of Black Hole Event Horizons:
• typical Galactic stellar mass Black Hole: 0.001 µ arcsec
• Galactic center Black Hole:
8
µ arcsec
• M31 Black Hole:
4
µ arcsec
Apparent Size of Black Hole Event Horizons:
• typical Galactic stellar mass Black Hole: 0.001 µ arcsec
• Galactic center Black Hole:
8
µ arcsec
• M31 Black Hole:
• M32 Black Hole:
• other nearby galaxies:
µ arcsec
0.07 µ arcsec
<1
µ arcsec
4
Other reasons why M31 is special:
• Next supermassive black hole beyond our Galaxy (+ M32)
• Black hole is ~30 times more massive than in the Galaxy
(see below) and large enough to produce a serious AGN.
• Chance to resolve single stars some day.
• Chance to exclude BH alternatives and study horizon.
• Very different nuclear structure from Galaxy…
• Not obscured by dust and gas.
Wendelstein-Calar Alto Pixellensing Project
16’x16’
30’’x30’’
NOAO
Wendelstein FWHM ~1.0’’
M31 center has more complex structure than the Galactic
center: it’s asymmetric!  is it in equilibrium?
Wendelstein 30”x30’’
FWHM ~ 1”
Real color:
F300W, F555W, F815W
Lauer et al. AJ 1998
Kormendy & Bender 1999
1” ~ 3.7 pc ~ 10 l.y.
M 31 with the Hubble Space Telescope
The nuclear structure of M31:
Real color:
F300W, F555W, F815W
Lauer et al. AJ 1998
Kormendy & Bender 1999
1” ~ 3.7 pc
P1
P2
P3
 • P1 and P2 are signatures of an eccentric disk of red stars,
not two independent nuclei (Tremaine 1995, Peiris & Tremaine
2003, Kormendy & Bender 1999, Bender et al. 2005). At a separation of
P1 and P2 of ~ 0.5” ~ 1.8pc and a rotation velocity of V ~ 200 km/s the
orbital period is only ~ 50000 years: two clusters would merge quickly.

• What is the nature of the blue nucleus P3? ( Part 2)
Peiris & Tremaine’s (2003) latest version of the eccentric disk model:
P1
P2
RB’s version of
the eccentric disk
P3
Peiris & Tremaine’s (2003) latest version of the eccentric disk model:
P1
P2
P3
An eccentric thin disk with inclination ~ 55o and an appropriate population
of orbits can reproduce both photometry and ground-based kinematics.
The kinematics require a black hole mass of: MBH ~ 108 M.
HST STIS spectroscopy in the NIR:
Bender, Kormendy, Bower et al. (2005): high
resolution STIS kinematics agree with Peiris
and Tremaine (2003) predictions very well!
P1
P2
Slit for NIR Ca triplet spectroscopy
with STIS G750M on F555W image
Peri-center disk stars show up in the STIS line-of-sight velocity distribution:
V ~ 1000 km/s is reached at r ~ 0.1”: This provides a simple mass estimate:
 M ~ V2r/G ~108 M
+
+
+
+
+
+
+
−
+
−
−
−
−
−
−
radius
scale
flipped
and
offset
by 0.05”
HOWEVER, one issue remains a problem:
Peiris & Tremaine assumed that the disk has no mass. In reality the disk
stars have ~107 M , i.e. their self-gravity cannot be neglected.
This creates non-closed orbits and orbit precession….
Bacon et al. 2001 suggest that P1+P2 can be explained by an m=1
mode in a disk. They study it with an N-body simulation. A pretty
convincing model with a pattern speed of ~ 2 million years is obtained.
However, it does not fit as well as the Peiris & Tremaine model.
face-on model
projected model
rotation curve
The nature of the blue nucleus P3
HST STIS spectroscopy in the blue and NUV
• mV = 18.7 (0.3)  MV = -5.7 (0.3)
• is it blue AGN continuum (King et al.)?
P3
• colors similar to A-type star (Lauer et al)
• if P3 is indeed composed of A-stars,
its light can come from:
 ~ 1 A0I (excluded), or
 ~ 200 A0V / A0III, or
 ~ 107.5 DA white dwarfs ????
Slit for blue nucleus spectrsocopy
with STIS G430L on F300W image
Hubble Space Telescope STIS spectroscopy of P3 ...
… reveals: The blue nucleus is NOT an AGN accretion disk
but a cluster of A-type stars.
on blue
nucleus
on background
P3
blue nucleus
Which type of A-stars constitute P3?
Did we
possibly
discover
a massive
cluster of DA
white dwarfs
?
Koester DA models
well, let’s play with the parameters T and log g …
NO! … this does not work, the P3 light is not dominated by DA’s
• Pure A0V, A0III, and A0I-populations as well as a 200 Myr
burst population fit the P3 spectrum equally well.
• A0I’s are excluded because one A0I star would already
have the luminosity of P3, but P3 is clearly extended.
• Stellar mergers cannot form A-stars from the red background
population, because multiple non-destructive mergers are
required and encounter speeds are too high (~1000 km/s).
A young
burst population
seems to provide
the most likely
explanation for the
spectrum of P3.
SSP models of
Bruzual& Charlot
yield:
200 Myr SSP:
M ~ 5000 Msun,
N ~ 15000 stars
 get back to this
later with ACS
imaging
Modeling photometry and kinematics of P3
real color
300 nm − eccentric disk
deconvolved
individual stars are starting to be resolved
300 nm
300 nm − eccentric disk
− P3 model
300 nm − eccentric disk
PSF convolved P3 model
Modeling photometry and kinematics of P3, continued…
300 nm − eccentric disk PSF convolved P3 model
P3 has an exponential profile, its
flattening is high and consistent
with a disk seen under the same
inclination as the larger P1+P2
eccentric disk: inclination ~ 55o!
 if P3 is really a disk,
then it should rotate fast …
P3: observed kinematics with HST STIS
a thin disk model for the kinematics of P3:
=> once the photometry is given, only the black hole mass
is a free parameter, and V and σ profiles should fit
surface brightness (PSF-c.)
V-field (PSF-convolved)
σ -field (PSF-convolved)
slit-pixels
0.5”x0.5”
0…1
-700 … +700 km/s
150 … 1000 km/s
P3 disk model (PSF-c.) V-field (PSF-convolved) σ -field (PSF-convolved)
P3: observed kinematics vs
thin disk model (+ Plummer MDO)
slit-pixels
0.5”x0.5”
0…1
-700 … +700 km/s
150 … 1000 km/s



• The best fit is obtained for a point
mass, i.e. a black hole of mass:

MBH ~ 1.4 x 108 M
• A Plummer sphere of radius 0.03”
is 1-sigma off from the BH solution.
b/a =0.26
The overall best fit to P3 is obtained for a thin exponential disk (i=55o),
Thick disk/spheroid, orbit-superposition models (higher inclin.) are less likely.
First HST WFPC imaging in UV, V, I
26
Real color WFPC image: F300W, F555W, F815W; Lauer et al. AJ 1998, Kormendy & Bender 1999
New HST ACS imaging in U and B
27
Bender, Lauer, Kormendy et al. 2008, real color UACS+BACS+IWFPC
HST ACS U
28
200 Myr pop,
IU - α IB , α = (fu/fb)P1,
i.e. P1-P2 disappear
29
U-B
M 31 P3 (U-band,
P1-P2 subtracted)
simulated images 100 Myr
simulated images 200 Myr
simulated images 400 Myr
30
power spectrum of surface brightness fluctuations
(mean exponential surface brightness model subtracted)
31
Is the MDO a black hole?
Is the MDO a black hole?
Is the MDO a black hole?
The 1-σ limit on the MDO size is rh = 0.11 pc.
But then M = 2.15 × 108 M to fit kinematics.
Ring Nebula (radius = 0.11 pc) at same scale
Astrophysical arguments rule out dark clusters:
1. Brown dwarfs would collide, get distroyed,
(or merge) and turn into gas (or visible stars)
Mbrown dwarf = 0.04
Mbrown dwarf = 0.01
Mbrown dwarf = 0.001
Mbrown dwarf = 0.0001
Mbrown dwarf = 0.00001
M ⇒ tcollision = 2.5 x 107 yr
M ⇒ tcollision = 4.4 x 106 yr
M ⇒ tcollision = 4.4 x 105 yr
M ⇒ tcollision = 9.7 x 104 yr
M ⇒ tcollision = 3.4 x 104 yr
If the MDO is made of any of the above brown dwarfs,
collisions easily convert them into gas in «1010 yr.
2. What about clusters of white dwarfs,
neutron stars or stellar mass black holes?
1 - Let’s form (M = 2.15 x 108 M, rh = 0.11 pc) via
n successive generations of progenitor stars,
where n = 1010 yr / (stellar lifetime).
2 - All stars of all generations have the same mass.
3 - Gas for star formation is provided.
4 - Don’t worry about the fact that mass loss during
stellar evolution unbinds the young MDO.
5 - Form stars with density proportional to that of MDO
(Plummer sphere).
6 - MDO has rh ∝ M-1. Compute evolution back in time from
the present.
This goes wrong for all types of remnants:
1 - MDO progenitors have MV = -16.3 to -17.5
throughout their formation. Such bright nuclei
could not be hidden in nearby galaxies. It is
unreasonable to assume that formation lasted for
1010 yr and stopped recently in all galaxies.
If formation took < 1010 yr, all problems get worse.
2 - Progenitors are more massive than remnants.
Dynamical friction deposits them at small radii.
Therefore it is impossible to make an MDO that is
as compact as a Plummer sphere.
3 - After a fraction of the MDO has been formed, the MDO
velocity dispersion is many times the surface escape
velocity of the newly formed stars. Collisions destroy stars.
If the MDO is less compact, then it must be heavier
and all problems get much worse.
This goes wrong for 0.6 - 1 M white dwarfs:
1 - Interior to rh, most progenitors collide and get destroyed.
If they succeed to merge, they get converted into progenitors
of 1 M white dwarfs, even neglecting dynamical friction.
Lower-mass white dwarfs are also irrelevant –
their progenitors get destroyed (or live more than 1010 yr
and therefore would be visible).
This goes wrong for ≥ 1.0 M white dwarfs:
1 - Including dynamical friction, most progenitors interior to rh
either get destroyed or converted into progenitors of highermass remnants.
2 - Except near the Chandrasekhar limit, white dwarfs at r ≤ r1/4
in the completed MDO collide and merge quickly enough so
that there should be a Type Ia supernova every ≤ 100 years.
They would easily be seen in distant galaxies.
This goes wrong for neutron stars and stellar BHs:
1 – If they don’t get destroyed, progenitors
interior to rh get converted into progenitors of
higher-mass remnants.
2 - NS and BH progenitors die in Type II supernovae.
During MDO formation, there should be a
supernova every ~ 100 – 200 years. They would
easily be seen in nearby and in distant galaxies.
3 - Gas expelled in supernovae easily escapes from
the MDO potential well ⇒ must provide new gas
for each generation of stars (~ 103 times).
Is the M 31 MDO a black hole or a cluster of dark objects?
half assembled
¾ assembled
fully assembled
applies to inner
¼ of mass.
clusters built
by maximum
number of
successive
generations
of stars!
Conclusions:
• the blue nucleus P3 in M31 is a thin to moderately thick stellar disk
in Keplerian rotation around a black hole.
• the P3 stellar population is dominated by A-type stars. It is very
unlikely that these can be formed in collisions of low mass stars.
• a 200 Million year old starburst population provides the best
explanation for P3’s spectrum and the power spectrum of P3’s
surface brightness fluctuations.
• P3 could consist of stars that formed in the outer parts of a dense
accretion disk; the gas mass needed is ~105 Msun.
+1.1
• MBH = 1.4 -0.3 108 M (1σ, allowing for different inclinations).
• astrophysical alternatives to a black hole in the form of dark
clusters can be ruled out because even under the most favorite
conditions they either evaporate or they cannot be built because
the progenitors of the compact objects destroy each other.
It’s a black hole!
Copyright NASA