Stellar Evolution
Stars are ~black bodies, i.e., in thermal
equilibrium.The radiation has a specific
spectrum and intensity that depends only on
the temperature of the body.
Stellar Evolution
More massive stars must have a higher temperature to maintain equilibrium.
Stars are approximately black-bodies, thus
hotter stars are also bluer.
Along the main sequence, temperature, T,
and luminosity, L, are related:
L
L
T
M
4
3.5
Stellar Evolution - Lifetimes
How does stellar mass relate to a star’s lifetime?
The total energy released by a star in its lifetime is, Etotal:
Etotal = L * time
Nuclear fusion is turning mass into energy, which means Etotal = eff * M c2
Using the relationship between L and M on the main sequence:
t
M
2.5
Massive star lifetimes are much shorter than lower mass stars.
Magnitude/Mass
Stellar Evolution
10-100 million years
(very short lifetime)
10Msun
2-8 billion years
Sun
0.1Msun
Color/Temperature
10-100 billion years
(longer than age of Universe)
Stellar Evolution
If stars of all masses form at same
time, can determine age
by noting which stars are just
evolving off of the main sequence.
Composite CMD for open
star clusters in the Milky Way
Science Interlude:
Color Magnitude Diagrams
HB =
Horizontal Branch
RGB = Red Giant Branch
MS =
main sequence
MS
WD = White Dwarfs
12 Billion years old
Chemical abundance similar to early universe
100 million years old
Chemical abundance similar to Sun.
Significant binary star fraction
Stellar Evolution:
Age and Metallicity
•
Open cluster members
have roughly the same
age and therefore
average metallicity.
What do astronomers consider a metal?
Astronomers consider any chemical element heavier than hydrogen a metal. This is because when
the Universe began in the Big Bang the only elements produced in large abundances where hydrogen and helium
(the two lightest elements). Stars produce all elements heavier than hydrogen in their cores as products of fusion.
What is metallicity?
Metallicity is a measure of the amount of metals in an object of study. Most often astronomers use measurements
of absorption lines in a stellar spectrum to measure the amount of iron present. Absorption lines due to hydrogen
are analyzed in a similar manner to obtain the amount of hydrogen. The ratio of the amount of iron to the
amount of hydrogen in the object is divided by the ratio of the amount of iron to the amount of
hydrogen in our Sun to obtain a metallicity relative to the Sun. This value, denoted as [Fe/H] for that object, is
plotted on a logarithmic scale.
•
[Fe/H] = -1 ---> 1/10th Solar amount.
•
[Fe/H] = +1 ---> 10x Solar amount.
https://edocs.uis.edu/jmart5/www/rrlyrae/metals.htm
What do we learn from a star's metallicity?
All of the iron and other metals in the universe have been produced in stars. At the end of a star's life, it recycles some
or all of the elements it has produced back into the interstellar medium. This processed material becomes mixed into
clouds where the next generation of stars are born. So each subsequent generation of stars is enriched with the metals
produced in previous generations. Because of this we can infer in a closed system, like our galaxy, that
stars with a lower metal content (smaller [Fe/H]) are older than stars with a higher metal content
(larger [Fe/H]).
Some context for our Galaxy:
Different parts of our galaxy have different metal contents so we can infer their relative ages. For example, the halo is
the part of the galaxy with the smallest amount of metals (the stars have an average [Fe/H] of -1.6 or about 30 times less
iron than the sun). For this reason we believe that the halo may be the oldest part of the galaxy. The lack of stars with
small metallicities in the disk of the galaxy lead us to believe that the disk was one of the last parts of the galaxy to
form.
Stellar Evolution:
Age and Metallicity
https://edocs.uis.edu/jmart5/www/rrlyrae/metals.htm
Image Credit: Swinburne University of Technology.
Stellar Evolution:
IMFs and Isochrones
•
Open cluster members
have roughly the same
age and metallicity.
•
An isochrone is a curve on the HR diagram, representing a
population of stars of a given age and metallicity.
•
If the Initial Mass Function (IMF) of a system of stars is
known, isochrones are calculated for a given age by taking
every star in the initial population and using numerical
simulations to evolve it forwards to the desired age.
Stellar Demographics
If a cluster has 5 stars at 10 Msun, then how
many does it have that are 1 Msun? 100 Msun?
Many more low-mass stars than high-mass stars form
Fig 16.20
Stellar Demographics
Stellar Initial Mass Function:
can get total stellar mass from tip of iceberg!
Fig 16.20
How big can stars get?
How can we find out?
Radiation Pressure
!Photons exert pressure when they strike matter
!Massive stars are so luminous that Radiation
Pressure drives massive stellar wind
Radiation Pressure wins:
Eddington Luminosity
The Bubble Nebula
NOAO
Stellar Maximum Mass
We're not sure…
! Radiation Pressure may limit max stellar mass
! Stars with > 150 MSun not seen
Max 150 MSun limited
Luminosity
by
radiation pressure?
Min 0.08 MSun
limited by fusion
capability
Temperature
Initial Mass Function
big vs small stars
Why does the IMF matter?
N(m) ~ m^(-a), where a const.
log N (log m)
log m
Cluster mass function is similar!
Initial Mass Function
The IMF specifies the fractional distribution in
mass of a newly formed stellar system. It is
often assumed to have a simple power law
(M) = c M
In general, (M) extends from a lower to an
upper cutoff, e.g., from 0.1 to 125 solar masses.
Commonly used IMFs are those of Salpeter
(1955), Scalo (1986), and Miller and Scalo
(1979).
http://webast.ast.obs-mip.fr/hyperz/hyperz_manual1/node7.html
Hillenbrand 1997, AJ, 113, 1733
Stellar Evolution:
IMFs and Isochrones
•
Open cluster members
have roughly the same
age and metallicity.
•
An isochrone is a curve on the HR diagram, representing a
population of stars of a given age and metallicity.
•
If the Initial Mass Function (IMF) of a system of stars is
known, isochrones are calculated for a given age by taking
every star in the initial population and using numerical
simulations to evolve it forwards to the desired age.
•
Isochrones as diagnostics: this theoretical isochrone can
be compared against the observed CMD to determine how
well they match. If they match well, the assumed age and
metallicity of the isochrone is close to that of the observed
system.
http://stellar.dartmouth.edu/models/isolf_new.html
Stellar Evolution:
IMFs and Isochrones
•
Hilker et al. (2004)
What do these
isochrones look like?
Investigators:
Investigator
Institution
Country
PI&
Dr. Nitya Kallivayalil
Yale University
USA/CT
CoI
Dr. Gurtina Besla
Columbia University in the City of New York
USA/NY
CoI
Dr. Roeland P. van der Marel
Space Telescope Science Institute
USA/MD
CoI
Dr. Szymon Kozlowski
Warsaw University
POL
CoI
Prof. Marla C. Geha
Yale University
USA/CT
CoI
Dr. C. S. Kochanek
The Ohio State University
USA/OH
CoI
Dr. Jay Anderson
Space Telescope Science Institute
USA/MD
Harvard University
USA/MA
CoI
Prof. Charles R. Alcock
Number of investigators: 8
& Phase I contacts: 1
Target Summary:
Target
RA
Dec
Magnitude
AGNC-S0162
00 37 4.6700
-73 22 29.60
V = 18.89 +/- 0.03, I=18.49
AGNC-S0186
00 38 57.5400
-74 10 0.90
V = 18.41 +/- 0.02, I=17.75
AGNC-S0199
00 39 47.8200
-74 34 44.80
V = 18.43 +/- 0.02, I=17.55
AGNC-S0202
00 39 57.6500
-73 06 3.60
V = 19.85 +/- 0.06, I=19.43
AGNC-S0246
00 42 59.0000
-74 02 44.60
V = 19.16 +/- 0.04, I=18.55
AGNC-S0281
00 44 40.2600
-73 21 51.80
V = 18.47 +/- 0.02, I=17.62
AGNC-S0291
00 45 7.5400
-72 41 21.60
V = 20.1 +/- 0.07, I=19.38
AGNC-S0295
00 45 16.7500
-74 42 31.10
V = 19.48 +/- 0.04, I=18.82
AGNC-S0659
00 54 22.9800
-73 31 0.20
V = 18.99 +/- 0.03, I=18.47
AGNC-S9999
00 56 13.2900
-72 38 20.60
V = 19.89 +/- 0.06, I=18.81
AGNC-S0826
01 00 5.7200
-71 57 23.40
V = 19.09 +/- 0.03, I=18.66
AGNC-S0834
01 00 18.2700
-74 03 22.80
V = 18.59 +/- 0.02, I=18.02
AGNC-S0867
01 01 4.7200
-73 41 59.90
V = 18.34 +/- 0.02, I=17.91
AGNC-S0993
01 02 44.9100
-72 15 21.90
V = 18.73 +/- 0.03, I=18.34
AGNC-S1072
01 05 22.5400
-71 56 49.90
V = 19.07 +/- 0.03, I=18.32
AGNC-S1120
01 07 15.6400
-74 10 45.30
V = 18.14 +/- 0.02, I=17.35
AGNC-S1124
01 07 21.6300
-72 48 45.60
V = 19.44 +/- 0.04, I=18.53
AGNC-S1145
01 08 25.4300
-73 43 17.30
V = 19.3 +/- 0.04, I=18.43
AGNC-S1148
01 08 34.8500
-71 19 15.50
V = 19.33 +/- 0.04, I=18.71
AGNC-S1188
01 11 3.0000
-72 20 36.20
V = 19.18 +/- 0.04, I=18.49
AGNC-S1271
01 14 45.3500
-71 53 40.80
V = 18.9 +/- 0.03, I=18.38
AGNC-S1280
01 15 18.7000
-73 23 54.60
V = 19.38 +/- 0.04, I=18.70
Scientific Justification
1 Background
At a distance of ∼ 50 kpc from the Galactic center, the Large and Small Magellanic Clouds
(LMC and SMC) are our closest example of an interacting pair of dwarf galaxies. Evidence of
their ongoing interaction is clearly illustrated by the pronounced bridge of HI gas connecting
them, known as the Magellanic Bridge (Putman et al. 2003). All theoretical models of the
Magellanic System predict that the Magellanic Bridge formed via a recent tidal encounter
between the LMC and SMC (e.g., Gardiner & Noguchi 1996). High-precision proper motions
(PMs) of the Clouds made by our group illustrate that such an encounter is unavoidable
(Kallivayalil et al. 2013; hereafter K13). These measurements were made using two epochs
of ACS High Resolution Camera (HRC) data in Cycles 11 and 13 and a third epoch of
WFC3/UVIS data in Cycle 17, centered on fields with background QSOs. Backward orbital
integration using these present day PMs necessarily imply that the Clouds were much closer
to each other 100-300 Myr ago than their current ∼ 20 kpc separation.
There are, however, significant uncertainties in the exact nature of LMC-SMC interactions
and their orbit about the Milky Way (MW), driven by the larger uncertainties in the SMC
PM (a factor of 3, ∼ 0.06 vs. ∼ 0.02 mas/yr for the LMC). The errors in the SMC PM are
dominated by our lack of knowledge of the internal kinematics and structure of the SMC, and
by the limited number of background QSOs with which to probe these internal kinematics
and calculate a center-of-mass (CM) motion (5 SMC QSOs vs. 21 LMC QSOs).
1.1 Project Goals
We propose here to use a large increase in background SMC QSOs, spanning a significantly
larger extent of the SMC (see Fig. 1) to investigate its internal structure and possible rotation,
and therefore to also vastly improve its CM PM determination. A long-term GO proposal
with a 2-year baseline, targeting 30 background QSOs will allow us to measure the SMC CM
PM to 0.01 mas/yr (∼ 3 km/s). The number of QSO targets has been chosen to ensure a
secure measurement of the in-plane rotation of the SMC (see Description of Observations).
The new QSOs were selected based on their mid-IR/optical colors (Koz!lowski & Kochanek
2009) and X-ray emission (Koz!lowski et al. 2010), and then spectroscopically confirmed
with 2DF/AAOMEGA (Koz!lowski et al. 2011, and 2013 in prep). While there are 212 new
SMC QSOs in all, we have selected 30 of the brightest (median V ∼ 19 mag) that are also
uniformly distributed behind the SMC. This will allow us to definitively constrain whether
the LMC & SMC are in a binary, whether they are on their first infall into the MW, and
the upper limit of the MW mass (Section 2), as well as the internal kinematics/rotation of
the SMC and whether it is a dwarf in transition (Section 3).
2 The Binarity of the LMC/SMC and the Orbital History of the Clouds
The third epoch PMs (K13) combined with recent revisions in the local standard of rest
velocity for the Solar motion (McMillan 2011; ∼ 240 km/s vs. the IAU standard of 220
km/s) and updated models for the internal rotation of the LMC, constrained by the PM
1
data itself (van der Marel & Kallivayalil 2013), have decreased the inferred 3D space motions
of the Clouds by ∼ 60 km/s (K13, Kallivayalil et al. 2006a,b). While the original velocities
strongly suggested that the Clouds are on their first infall to our system (Besla et al. 2007),
the reduced velocities have reopened debate concerning the accretion epoch. The velocities
are still too high for short-period (< 2 Gyr) orbits to be viable; however, it is possible that
the Clouds have completed one orbit within the past 6 Gyr (e.g., Shattow & Loeb 2007, Diaz
& Bekki 2012) if the MW is sufficiently massive (> 1.5 × 1012 M⊙ ). At the same time, a first
infall scenario is also not ruled out (K13).
Furthermore, the observed relative velocity between the Clouds (128 ± 32 km/s) is of order
the escape speed of the SMC from the LMC, making it difficult to maintain a long-lived
binary state while the Clouds are subject to the MW’s tidal field. In K13 we explored the
error space and found that the past orbits of the Clouds about the MW and their own
interaction history are connected, such that the key discriminant between a first infall or
early accretion event is the assumption that the SMC has existed in a long-lived (>
∼ 4 Gyr)
binary configuration with the LMC. As the mass of the MW increases, it is more likely that
the LMC has completed an orbit about the MW. However, at the same time the MW’s tidal
field is increasingly efficient at disrupting the LMC-SMC binary. Specifically, we find that
long-lived binary states are preferentially found in first infall orbits (see Fig. 2).
However, the error bars on the relative motion are still quite large, allowing for multiple
solutions and degeneracy with the mass of the LMC and MW. Smaller SMC PM errors (of
the order of the LMC PMs) are required to uniquely determine whether the Clouds are
indeed bound to each other, providing tighter constraints on their interaction history with
the MW.
2.1 The Mass of the Milky Way
Given the SMC errors, there is still a small probability of long-lived binary configurations
in higher mass MW models, which might allow for a previous passage of the Clouds about
the MW. If the PM error space of the SMC were reduced, the K13 model can be used to
place strong upper bounds on the mass of the MW. Specifically, if the mass of the MW
is too large, stable binary configurations may be impossible for reasonable estimates of
the LMC’s gravitational binding energy, given the SMC’s constrained velocity error-space.
While satellite orbits can place lower bounds on the halo mass of the MW (e.g., the space
motion of Leo I; Sohn, Besla et al. 2013), the upper limit for the MW halo mass is largely
unconstrained, apart from the Local Group timing argument (van der Marel et al. 2012).
3 Internal Kinematics of the SMC
The Clouds are often referred to as ‘Irregular’ type galaxies, where the term ‘Irregular’
implies a lack of organized structure (Wilcots 2009). Upon closer inspection, the SMC is
not quite as ‘irregular’ as it might originally appear. Its older stellar distribution is better
described as a dispersion-supported spheroid (Harris & Zaritsky 2006), meaning its irregular
appearance in the optical is largely due to disorganized, on-going star formation. Given the
surprising lack of irregularity exhibited by the Clouds, De Vaucouleurs & Freeman (1972)
2
concluded that the MW had very little to do with their evolution, in sharp contrast to the
prevailing view at the time that the Clouds were long term companions to the MW. This
means that tidal interactions between the Clouds themselves become more important in
their morphological evolution. Critical to understanding the strength of such interactions
are accurate measurements of their internal kinematics.
We were able to use the relatively large number of LMC QSOs, and the excellent quality of
the data, to derive a purely PM-based rotation curve of the LMC (K13). Unlike line-of-sight
(LOS) velocity studies, the PM field can constrain the inclination of the disk independently.
In Fig. 3 (right; van der Marel & Kallivayalil 2013) we show the three-epoch rotation curve
of the LMC. The weighted average of the outer-most data points gives a value of 76 ± 7
km s−1 , which is in excellent agreement with the latest LOS velocity-based studies (Olsen
et al. 2011), as well as the HI value (Kim et al 1998). In this proposal, we would like to
investigate SMC rotation in the plane of the sky, something that we were unable to constrain
in K13 with so few QSOs.
In contrast to the LMC, relatively few kinematic tracers have been studied for the SMC (van
der Marel et al. 2009). While the older stellar population exhibits little to no LOS rotation
in the SMC, the gas and young stars show a pronounced velocity gradient. The rotation
curve of the HI gas peaks at ∼ 60 km/s (0.2 mas/yr) at a radius of 3 kpc (∼ 3◦ ; Stanimirović
et al. 2004). Evans & Howarth (2008) have found a velocity gradient of similar slope in the
kinematics of young (O,B,A) stars. This gradient is oriented almost orthogonally to that of
the HI. Van der Marel et al. (2009) note that this may be a result of the different spatial
coverage of the two studies, as the Evans & Howarth study did not cover the North-East
region of the SMC. This will be verified with our PM program.
Harris & Zaritsky (2006) examined the older RGB population within the inner few degrees
of the SMC and found σ = 27.5 ± 0.5 km/s with little rotation. The low Vrot /σ is consistent
with that of dE and dSph galaxies, suggesting that the SMC may be a transition galaxy
between gas rich, rotation supported dIrrs and dispersion supported dSphs. However, it is
known that the SMC has a considerable LOS depth, and it may be viewed nearly pole-on
(e.g., Crowl et al. 2001). So it may have significant in-plane rotation that is not manifest
along the LOS. Our study will test this.
3.1 Is the SMC a Dwarf in Transition?
We have recently put forth a theory that explains this kinematic mismatch and identifies a
possible transition mode between dIrr and dSph galaxies: that of a direct collision (Besla
et al. 2012). The SMC is initially modeled with a well-defined stellar and gaseous disk, but
undergoes a direct collision with the LMC ∼ 100 − 300 Myr ago. During this dramatic
encounter SMC stars are subject to shocks and tidal heating, which wipe out the initial
velocity gradient of the SMC’s stellar disk, stripping stars and gas from the inner regions of
the SMC to much larger radii. The gas, however, is dissipative and is thus able to cool and
retain the velocity gradient of the original disk (see Fig. 4). Our PM program can validate
this theoretical model for the internal kinematics of the SMC.
3
References
Anderson, J. & King, I. 2006, ACS Instr. Science Rep. 06-01
Anderson, J., & van der Marel, R. P. 2010, ApJ, 710, 1032
Besla, G., et al. 2007, ApJ, 668, 949
Besla, G., et al. 2010, ApJ, 721, L97
Besla, G., et al. 2012, MNRAS, 421, 2109
Crowl, H. et al. 2001, AJ, 122, 220
de Vaucouleurs & Freeman 1972, Vistas in Astr., 14, 163
Diaz, J. D. & Bekki, K. 2012, ApJ, 750, 36
Evans, C. J., & Howarth, I. D. 2008, MNRAS, 386, 826
Gardiner & Noguchi 1996, MNRAS, 278, 191
Geha, M., et al. 2003, AJ, 125, 1
Kallivayalil, N., et al. 2006a, ApJ, 638, 772
Kallivayalil, N., et al. 2006b, ApJ, 652, 1213
Kallivayalil, N., et al. 2013, ApJ, 764, 161
Kim, S., et al. 1998, ApJ, 503, 674
Koz!lowski, S., & Kochanek, C. S. 2009, ApJ, 701, 508
Koz!lowski, S., et al. 2010, ApJ, 708, 927
Koz!lowski, S. et al. 2011, ApJS, 194, 22
McMillan, P. J. 2011, MNRAS, 414, 2446
Olsen, K. A. G., et al. 2011 ApJ, 737, 29
Putman, M. E., et al. 2003, ApJ, 586, 170
Shattow, G. & Loeb, A. 2009, MNRAS, 392, L21
Sohn, S. T., et al. 2012, arXiv:1210.6039
Sohn, S. T., et al. 2012, ApJ, 753, 7
Stanimirović, S., et al. 2004, ApJ, 604, 176
van der Marel, R. P. et al. IAU Symposium, Vol. 256, p. 81
van der Marel, R. P., et al. 2012, ApJ, 753, 8
van der Marel, R. P., & Anderson, J. 2010, ApJ, 710, 1063
van der Marel, R. P. & Kallivayalil, N. 2013, ApJ-submitted
Wilcots, E. M. 2009, IAU Symposium, 256, 461
Figure 1: R-band image of the SMC (3◦ × 5◦ ). The MACHO photometric coverage is indicated.
Red symbols indicate our new reference QSOs spectroscopically identified with AAOMEGA
(Koz!lowski et al. 2011; and in prep). The yellow symbols show the old QSO sample used in K13
and Kallivayalil et al. (2006b).
Description of the Observations
1 Observational strategy
We propose to observe 30 quasars, obtaining two epochs of WFC3/UVIS data over two cycles
separated by the maximum time interval allowed by scheduling. Each epoch will consist of
a one orbit observation, obtaining 4 dithered (BOX pattern to optimize PSF sampling)
exposures at V (F606W) and two, shorter, dithered exposures at I (F814W). The QSO
brightness ranges from 17.7 ≤ V ≤ 20.1, with a median brightness of V = 19.1. We aim for
4
Figure 2: Left: An orbit in which the LMC (solid red line) and the SMC (dashed red line) are
bound for the past 4 Gyr, given the mean K13 LMC velocity, and from searching the error space
for the SMC velocity (and Mvir = 1.5 × 1012 M⊙ , MLMC = 1.8 × 1011 M⊙ ). Right: These same
parameters give a Galactocentric orbit in which the Clouds (LMC in red, SMC in green) are on
their first infall into the MW. For reference, the blue line shows the relative distance between the
Clouds. The dot-dash horizontal line shows the virial radius, Rvir , for this MW model.
a S/N > 200 for the quasars in F606W, and we need the F814W data to make CMDs to
separate the SMC from field stars. We also request coordinated parallel observations with
ACS in F606W and F814W. Since these parallel fields will fall on random SMC fields, we
will use the additional coverage to learn more about the stellar populations of the SMC. Our
previous work has shown that a two-year baseline is sufficient to achieve our target accuracy
(see Kallivayalil et al. 2006a,b, and below) and we therefore apply for time in this cycle and
in Cycle 23.
2 Why QSOs?
Recent proper motion studies have shown that accurate PMs can be achieved using background galaxies rather than QSOs (Sohn et al. 2012). While an individual background galaxy
is not as well measured as
√ a point source (QSO), the fact that there are many background
galaxies can give a large N advantage in the final accuracy. However, background galaxies
require deeper exposures, and therefore fewer individual pointings are possible, which would
not be well-suited for mapping the SMC velocity field. Also, background galaxies cannot be
used in crowded star fields (for e.g., near the SMC center). Therefore we think that QSOs
are more appropriate for this purpose, although if any relatively isolated background galaxies
are found in our images we will use them in our analysis. Given the FOV we do not expect
multiple QSOs per field, but if present, they would also allow additional consistency checks
to each field measurement.
5
Figure 3: Results from our previous work in K13 and van der Marel & Kallivayalil (2013). (Left)
The x vs. y (equivalent to W, N ) positions in pixels of the QSO, after transformation into the
masterframe (in which stars are centered around zero by construction), for 1 QSO field, for all 3
epochs of data. Each triangle represents 1 of 20 positions of the QSO relative to the starfield.
The uncertainty in centering the QSO within a given epoch is visible as scatter between the
points. The QSO’s reflex motion with respect to the stars is evident. (Right) The LMC PM
rotation curve, transformed to km/s, as function of cylindrical radius in the disk (in kpc). Thin
black points are from our two epoch ACS reanalysis. Red thick points are from the three-epoch
PM data. These data are well fit by the line-of-sight rotation curve of many tracers inferred
independently by Olsen et al. (2011; blue dot-dash line). The very good agreement between
line-of-sight and PM rotation measures of the LMC provides additional evidence for the accuracy
of the PM analysis.
3 Astrometric Precision
We know from our previous work (K13) that there are many SMC stars in each UVIS field,
roughly 800 well-measured stars versus ∼ 30 well-measured stars for the HRC. This will
significantly aid us in our aims here. We expect to be able to centroid similarly as in our
HRC data. Even though the WFC3 pixels are somewhat bigger, with adequate dithering this
does not significantly degrade the astrometry because the astrometric error is proportional
to the (FWHM/SNR) of the target. In fact, the UVIS often does better than the HRC (see
Fig. 3 (left)). In Fig. 3 (right), we show our measurement of the LMC rotation purely from
our PM data (van der Marel & Kallivayalil, 2013) as proof that our method works.
The typical QSO position error in our HRC data, with the 2-yr baseline (the same as
what has been proposed here) was 0.07 mas/yr = 20 km/s. By averaging√many fields, the
in-plane rotation velocity, Vrot , of the SMC will be determined to a factor N better, where
N is the number of fields. We wish to determine Vrot to ∼ 20% accuracy, which dictates that
N = 30, and this is why we request 30 QSO-fields.
We also found that individual stars of similar brightness as the QSO could be measured
6
Figure 4: The kinematics of gas and old stars from Model 2 of Besla et al. (2012), wherein the
SMC collided directly with the LMC (impact parameter < 5 kpc) within the past 100-300 Myr.
The kinematics of the SMC gaseous disk (left), and old stars (right) within 3 kpc (∼ 3◦ of the
SMC, illustrated in the LOS frame. There is a well-defined velocity gradient of 40 km/s in the
gas, whereas the gradient is < 10 km/s in the old stellar component. While the magnitude of the
HI gradient is lower than that observed (Stanimirović et al. 2004), the overall direction of the
gradient is correct. Despite modeling the SMC with a well-defined stellar disk initially, after the
collision the old stars do not display the clear gradient seen in the gas and young stars. There
may be a small gradient, but it is no larger than 10 km/s, consistent with the results of Harris &
Zaritsky (2006).
at this precision. Therefore, the overall rotation will be very well measured since that can be
averaged over stars. But differential rotation measurements between different stellar populations will be attainable as well. For comparison, the expected rotation signals of various
populations are in the 0.1–0.2 mas/yr range (the latter if there is a stellar rotation component as high as the HI). Measurements of the internal velocity dispersions for populations of
RR-Lyrae will be very well-measured, as well as for populations with very distinct kinematics
from the SMC, for e.g., an SMC population accreted onto the LMC (see Olsen et al. 2011).
Finally, with 30 QSOs for the SMC this will result in CM PM random errors of ∼ 0.01
mas/yr per proper motion component, which is of the order of but better than we have done
for the LMC. However, as we have argued, the SMC PM is by far the rate-limiting step in
our full characterization of the Magellanic system and its interaction with the MW. The key
gain here is the number and distribution of the QSOs which will allow us to better constrain
the SMC geometry and directly probe internal kinematics, and therefore also allow a secure
CM PM measurement.
4 Team Expertise
Our team has been critically involved in establishing HST as a key astrometric instrument.
Anderson & King (e.g., 2006) have performed the geometric distortion and PSF callibrations.
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These have been used by our Co-Is to derive the most accurate proper motion (PM) catalog
of any globular cluster: for Omega Cen 170,000 PMs with 70µas/yr median per star accuracy
(1.6 km/s; Anderson & van der Marel 2010) have been used to constrain the presence of an
intermediate-mass black hole (van der Marel & Anderson 2010). We also produced the most
accurate PM determinations to date for the LMC and SMC (20 and 60µas/yr; Kallivayalil
et al. 2006a,b,2013), and measured the first PM rotation curve for any galaxy (Fig. 2; van
der Marel & Kallivayalil 2013). We have put forth the theory that the Clouds may be
on their first passage, and that their own interactions far from the MW virial radius have
been the main driver of their evolution (Besla et al. 2007,2010,2012). Geha, Koz!lowski and
Kochanek have all been instrumental in finding the background QSOs and are experts in
this field (Geha et al. 2003, Koz!lowski & Kochanek 2009, Koz!lowski et al. 2011). Alcock has
been critically involved since the inception of this project and provides both theoretical and
observational input.
Special Requirements
The first and second epochs of this program will need to be obtained at the same ORIENT.
Coordinated Observations
Justify Duplications
We are requesting long-term GO observations of SMC fields, in order to determine proper
motions. The success of our project thus requires us to re-image these fields in the future.
Past HST Usage
Note that the description of past HST usage DOES NOT count against the page limits of
the proposal.
PI Summary: In the past four cycles, Kallivayalil has been PI on the HST Program
GO-11730, Cycle 17, “Continued Proper Motions of the Magellanic Clouds”. Data status:
completed.
The analysis of the Cycle 17 data is complete and published in Kallivayalil et al. 2013,
and the LMC rotation curve is submitted (van der Marel & Kallivayalil 2013). This program
was intended to refine and check the results from programs executed in Cycles 11 and 13
(PI: Alcock; Kallivayalil, Geha & van der Marel were Co-Is). As a group, data from these
programs have been published in several observational works as well as informed many
theoretical works that the PI has been involved in:
(1) Kallivayalil, N., van der Marel, R. P., Alcock, C., Axelrod, T., Cook, K. H., Drake, A. J.,
& Geha, M. 2006, ApJ, 638, 772
(2) Kallivayalil, N., van der Marel, R. P., & Alcock, C. 2006, ApJ, 652, 1213
(3) Besla, G., Kallivayalil, N., Hernquist, L., Robertson, B., Cox, T. J., van der Marel, R. P.,
& Alcock, C. 2007, ApJ, 668, 949
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Class Exercise: Evaluate my NOAO
Proposal
•
Some guiding questions:
•
Is the “big picture” question clear and well-framed? Or is it lost in
the details?
•
Is the sample (or target) justified? Why this particular target as
opposed to others?
•
Is this a timely investigation? Is the significance to astronomy of
the proposed program made clearly? Is there a clear discussion of
how it will further our understanding of an outstanding issue/
question?
•
Is the request for time (or desired depth of observations) justified?
•
Is the request for this particular Telescope justified? (e.g., FOV,
pixel scale or resolution, efficiency). Could the goals be better
achieved with another facility?