P1246990954Rukvr

1
University of Toronto 3 October 2003
New Features in the Inner and Outer Milky Way
SGPS team Naomi McClure-Griffiths, ATNF
Bryan Gaensler, Harvard Uni.
Anne Green, Sydney Uni.
Marijke Haverkorn, Harvard Uni.
Dain Kavars, Uni. Minnesota
Simon Strasser, Uni. Minnesota
Yosa Gelfand, Harvard Uni.
2
DRAO
NRAO GBT
The International Galactic Plane Survey
NRAO VLA
ATNF Compact Array
ATNF Parkes
3
The SGPS combines data from both
The Australia Telescope
Compact Array
The Parkes Telescope
4
Combining data from the single dish and the
interferometer
Vela supernova remnant with nearby HII region
5
Outline
Or as far as I get in 45 minutes.
6
21-cm line formulae The emission
coefficient
density n
erg
-33
j 1.6 10 n f( )
n
sec Hz sterrad
n
The brightness temperature
T ( ) j ds
n
n
B
-2
cm
18
-1
K km s
if the line is optically thin all along the line
of sight !
7

• We map the 21-cm line and continuum
• radiation in the longitude range 2550 - 3570,
• latitude -10 to 10 with resolution 2 arcmin
• latitude -100 to 100 with resolution 15 arcmin.

8
The Longitude-Velocity Diagram
9
Another longitude velocity diagram, showing the
outer galaxy velocities more clearly.
McClure-Griffiths et al., 2004, in prep.
10
Dotted line spiral connecting to the Norma Arm ?
solid lines Rg 16, 18, 20, 22, 24 kpc
11
(No Transcript)
12
(No Transcript)
13
Longitude-Latitude profile of the peak brightness
temperatures in the radial range 19 kpc. The warp causes the skew pattern, but the
more distant gas (Rg 25 kpc) follows b0 more
closely than the gas at Rg 20 kpc.
14
A face-on view of the location of the 10K
Tb threshold crossing.
15
Locus of Tangent Points
To study the inner Galaxy, we look at the...
At the tangent point
d
Rg
Nominal spiral arms
Model electron density Cordes Lazio (2001)
16
A deeper version of the l-v diagram ...
The terminal velocity edge is quite sharp.
17
Consecutive l-v diagrams averaged over 0.1 degree
steps in latitude from b 0.5o to b -0.5o.
18
threshold in TB
l325.5 deg
19
The threshold-crossing velocity as a function of
longitude and latitude
20
(aside)
Sunset over the ocean ? No, its just the
terminal velocity vs. latitude and longitude...
21
At each longitude, we look at the threshold
crossings at for every latitude, and see how they
are distributed in velocity.
We look at the distribution, find the median and
percentiles.
22
The Milky Way Rotation Curve Percentiles
(20, 50, and 80) of latitude slices within b
23
Northern () and Southern (.) Hemisphere data on
the terminal velocity vs. sin(l) RG / Rsun .
24
The Rotation Curve Depends on Qsun
km s -1
Rgal / Rsun
25
Terminal velocity data translated into
circular rotation velocity (the rotation curve)
26
Terminal velocity data translated into
circular rotation velocity (the rotation curve)
Brand and Blitz (1993)
27
Terminal velocity data translated into
circular rotation velocity (the rotation curve)
28
Terminal velocity data translated into
circular rotation velocity (the rotation curve)
Clemens 1985
29
Terminal velocity data translated into
circular rotation velocity (the rotation curve)
30
Terminal velocity data translated into
circular rotation velocity (the rotation curve)
Best fit Brandt Curve
V max 228 km s -1, R max 8.5 kpc, n 0.7
31
Terminal velocity data translated into
circular rotation velocity (the rotation curve)
32
Scutum-Crux
Norma
Sgr-Carina
Scutum-Crux
33
The Cordes and Lazio model (and the Taylor and
Cordes model) shows an arbitrary pitch angle
change in the Sgr-Car arm just inside the solar
circle. Could this explain the offset ?
Sagittarius - Carina Arm Pitch angle
variation ?
34
The density-wave model predicts that the gas
should slow down its rotational motion as it
approaches a spiral arm (from the concave side,
inside corotation) because the acceleration is
mainly radial, causing the gas to slow down to
conserve angular momentum. This is exactly what
we see !
35
Terminal velocity data translated into
circular rotation velocity (the rotation curve)
36
HI brightness temperature (density) at the
terminal velocity...
37
Measured at the terminal velocity (i.e. at the
sub- central point)
21-cm Brightness Temperature
Profiles in latitude look like
Galactic Latitude
Galactic Longitude
38
HI Disk Midpoint (past studies)
From Malhotra (1995) using published Parkes data
(boxes) and Dame et al (1987) CO data (line)
39
HI Disk Midpoint SGPS
Norma spiral arm
Scutum-Crux arm
3 kpc arm
40
HI Disk Scale Height (earlier studies)
From Malhotra (1995) using published Parkes data
41
HI Disk Scale Height (SGPS)
Scutum-Crux arm
3 kpc arm
Norma arm
flare?
42
The H I Disk Scale Height

• The scale height increases at spiral arms
• 3 kpc arm, Norma arm, Scutum-Crux arm
• What causes the increase?
• Increased turbulent pressure in areas of active
  star formation?
• Spiral density wave?
• The flaring of the disk for Rg strong, we see only an increase in scale height
  of 40 pc over 4 kpc radial range.

43
How does the brightness temperature (HI column
density) drop off beyond the terminal velocity ?
s v 7.6 km/s
An average of all spectra with 290 to align the terminal velocities at v0.
s v 16.8 km/s
44
On scales of a few hundred parsecs there are
different kinds of structure in the ISM. The H I
is everywhere, so the 21-cm line is good for
tracing the really BIG things. One very
common kind of structure is giant shells and
chimneys, presumably filled with hot, ionized gas
(HIM), perhaps blown by the collective action
of stellar winds and supernova explosions of many
stars.
45
A nearby interstellar shell similar to the local
bubble from the SGPS (Southern Galactic Plane
Survey)
Dickey 2001, ASP Conf. Ser. 231, p. 318.
46
GSH 2770036
McClure-Griffiths et al. 2003, Ap. J. submitted.
McClure-Griffiths, et al. 2000 A.J. 119,
2828. Parkes plus ATCA (1020 field) mosaic.
An HI bubble that blows out of the disk into the
halo above and below.
47
Parkes data only, velocities from 10 to 70 km/s.
48
Schematic face-on view
longitude-velocity (like a top view)
49
(No Transcript)
50
(No Transcript)
51
Some structures in the shell walls are easy
to understand, like these Rayleigh-Taylor fingers.
Others are more problematic, like these narrow
ridges that seem to join at the bottom.
52
GSH 2770036 is large, but not atypical of
supershells and chimneys in the outer Milky Way.
It appears to break out of the disk both above
and below.
53
On even smaller scales (tens of parsecs and
smaller) the structure of the ISM seems to be
stochastic (random). How are we to understand
this ?
54
(No Transcript)
55
(No Transcript)
56
Even noise has a spectrum. Here we compute the
spatial power spectrum of the emission.
Region 1. ATCA plus Parkes
Dickey, McClure-Griffiths, Stanimirovic,
Gaensler, Green 2001 Ap. J. 561, 264
57
Region 2. ATCA only (but mosaicked !)
58
The spatial power spectrum is the Fourier
Transform of the Structure Function
sky brightness distribution
fringe visibility function
Fourier conjugates
magnitude squared
autocorrelation
spatial power spectrum
structure function
Fourier conjugates
59
The Structure Function
In the ionized gas (the WIM) we see a turbulence
spectrum over many orders of magnitude in scale
size.
Figure from Armstrong, Rickett, and Spangler
1995, Ap. J. 443, 209.
60
Region 1. Same data, but now using channel
width 20 km/s instead of 0.82 km/s.
61
The slope of the power law changes with velocity
width !
62
Turbulence theory predicts that the slope should
steepen by one unit when we go from a thin slice
to a thick slice of the medium (Lazarian and
Pogosyan 2000 Ap. J. 537, 720).
63
The spatial power spectrum of a spectral line
tracer gives us the ability to trace the small
scale structure of the ISM dynamically. The
slope change with velocity width (depth) is
strong confirmation that the power law structure
function is actually tracing turbulence, rather
than some other random process with power law
statistics.
64
Next up Using absorption and emission
in the 21-cm line together to study the cool
phase vs. warm phase of the neutral atomic
hydrogen.
65
density n, temperature T
the optical depth
the equivalent width (velocity integral of the
optical depth)
EW
66
Wolfire, M.G., Hollenbach, D., McKee, C.F.,
Tielens, A.G.G.M., and Bakes, E.L.O., 1995, Ap.
J. 443, 152.
67
The shape of the cooling curve determines the
heating-cooling equilibrium values of density
and temperature.
Dalgarno and McCray, 1972 ARAA 10, 375.
68
Wolfire, M.G., Hollenbach, D., McKee, C.F.,
Tielens, A.G.G.M., and Bakes, E.L.O., 1995, Ap.
J. 443, 152.
pressure
density
69
McKee and Ostriker 1977, Ap. J. 218, 148.
Result the ISM pressure is always bouncing
around.
70
To study the cool phase gas, we look at 21-cm
absorption toward continuum sources.
But first what do the continuum sources in the
SGPS look like ?
71
high contrast
low contrast
SGPS mosaic of the 21-cm continuum emission in
the fourth Galactic quadrant.
72
There are many continuum sources in the SGPS
region. We filter the data to remove the
emission in the direction of the continuum, then
we measure the absorption and emission spectrum
pairs, and we get ...
73
(No Transcript)
74
What do we do with these ? 1. Study
the distribution of CNM in the Galaxy (for
this we use only the absorption spectra) 2.
Study the temperature of the CNM (for this
we combine the emission and absorption)
75
We compute the opacity, from the optical
depth integral divided by the corresponding path
length
EW is the equivalent width
is the line-of-sight averaged opacity
76
The galactic rotation curve tells us the velocity
as a function of distance along the line of
sight, v(s).
The velocity gradient, dv/dr, tells us the
path length corresponding to a given
bandwidth (e.g. one channel).
77
The velocity field due to Galactic rotation
sets v(s) , the radial velocity as a function of
distance along the line of sight.
s
Ds
v
DV
78
(No Transcript)
79
The opacity is larger in the inner Galaxy than it
is at the solar circle. It may be modulated by
the spiral arms.
80
The higher opacity in the inner Galaxy suggests
that either there is more cool phase gas
(relative to warm) or the median cool
phase temperature (Tcool ) is colder in the inner
Milky Way than at the solar circle, or some of
each.
This argues against most of the H I being
recently photo-dissociated H2.
e.g. Allen 2001
81
What is the temperature of the cool HI (CNM) ?
We can measure this by combining the emission and
absorption spectra channel by channel. But we
must somehow separate the emission from the WNM,
that shows no absorption because it is too warm.
82
If all the gas at a given velocity were at the
same temperature, we could measure the
excitation temperature (spin temperature)
directly
But in each velocity channel there is overlap of
several regions with different temperatures
83
We make the two-phase assumption, that
the absorption comes from the CNM only.
b
f
density n, temperature Tcool
The emission comes from the CNM, plus the WNM
(with brightness temp., Tw, proportional to
its column density) in front (f) and behind (b)
the cool gas.
84
What we just had was
Next subtract the continuum to get the line
brightness temp
Define x ( 1 - e-t )
The warm gas has much broader linewidths than the
CNM, so over the velocity range covered by one
absorption line we can approximate the WNM
emission by a linear function of v
85
Tw,b
Now define e to be the background fraction of Tw
, i.e. e
Tw,b Tw,f
We measure both x and TB for all velocities,
v, across each absorption line. We assume a
value for e and then least-squares fit the values
of co , c1 , and c2 separately for each
absorption line.
This is a linear least-squares fit. No first
guess, no local minima the solution is unique.
86
Example Fitting two absorption lines toward a
continuum source at intermediate latitudes (data
from Heiles and Troland, 2003, Ap. J. in press,
from Arecibo).
87
(No Transcript)
88
HISA HI self-absorption
Whenever the coefficient c2 is less than zero,
then the absorption line decreases the brightness
temp, i.e. this is HISA. Two out of three
absorption lines show c2
89
In the end, we get a distribution of CNM
temperatures among the different absorption
lines, depending on the assumed value of e.
The median value is 65 K, but note the tail
toward low temperatures (
90
What is the Distribution of HI Temperatures ?
The Old Picture
Log HI Column Density
1 2 3
4
Log Temperature (K)
91
What is the Distribution of HI Temperatures ?
WIM
The New Picture
Molecular clouds
WNM
PDRs
Log HI Column Density
Diffuse clouds
1 2 3
4
Log Temperature (K)
92
Different galaxies have different peak brightness
temperatures and different relative abundances
of warm and cool phase H I. This suggests
variations in some or all of Tcool , ,
nw , and our vantage point (face-on vs. edge-on).
In the SMC 21-cm absorption is rare, the ratio of
absorption to emission is much less than in the
Milky Way.
93
Dickey, J.M., Mebold, U., Stanimirovic, S., and
Staveley-Smith, L., 2000, Ap. J. 536, 756
94
Wolfire, M.G., Hollenbach, D., McKee, C.F.,
Tielens, A.G.G.M., and Bakes, E.L.O., 1995, Ap.
J. 443, 152.
Heating-cooling equilibrium is a strong function
of metallicity, z, and dust to gas ratio, D/G.
95
Conclusions

• From the SGPS data, we find
• the rotation curve of the H I
• the disk thickness, corrugations, warp and
  flare
• the giant shells through which processed
  material is returned to the cool phase
• the spatial power spectrum slope change
• the radial variation of the CNM abundance
• the distribution of CNM temperatures

96
The Cycle of Galactic Evolution
SN explosions
winds
nucleosynthesis
Star formation
97
l21-cm Galactic Plane Surveys trace...
SN explosions
winds
Turbulence
HI emission mapping
nucleosynthesis
HI absorption
Star formation
98
SGPS web site
http//www.astro.umn.edu/naomi/sgps.html
99
(No Transcript)
100

• Velocity integrals (over one channel, the entire
• spectrum, or any velocity range in between) give
• the column density, N

• for the emission spectrum (for low optical depth,
  t),
• the equivalent width, EW, for the absorption
  spectrum

Combining the two gives the density weighted
excitation temperature, Tspin . Tspin is
generally equal to the kinetic temperature in
the neutral gas.
101
Note that is not a physical temperature,
but a blend of warm and cool phases along
the line of sight. If we know the mean cool
phase temperature, Tcool , then tells us
the fraction of gas in the cool phase, fc
fc
Ncool
Tcool

Nwarm Ncool

The two-phase model comes from heating-cooling
equilibrium (Field, Goldsmith, and Habing, 1969,
Ap. J. Lett. 155, L149).
102
(No Transcript)
103
(No Transcript)