Infrared and Sub-millimeter
Astronomy
Introduction & Overview
Chris O’Dea
Acknowledgements: Steve Beckwith, Don Figer,
Bernie Rauscher, & Jeff Valenti
Outline
Historical Overview
 IR Detectors
 Backgrounds



The atmosphere
Astronomical
Radiative Processes
 IR Sub-mm Science
 NGST

What constitutes infrared ?

Traditionally 1 mm – 1 mm



Now: 1 mm – 300 mm




1 mm is long wave cutoff of silicon CCDs and
photographic emulsions
Initial mm-wave observations with bolometers
CCDs still limited; InSb/HgCdTe to ~0.6 mm
High frequency heterodyne receivers to <350 mm
Bolometers still dominate broad band to ~1.3 mm
Note: 1000 mm = 1 mm = 300 GHz
History








Herschel’s detection of IR from Sun in 1800
Johnson’s IR photometry of stars (PbS) mid 60’s
Neugebauer & Leighton: 2mm Sky Survey (PbS), late 60’s
Development of bolometer (Low) late 60’s
Development of InSb (mainly military) early 70’s
IRAS 1983
Panoramic arrays (InSb, HgCdTe, Si:As IBCs) mid-80’s
NICMOS, 2MASS, IRTF, UKIRT, KAO, common-user
instruments, Gemini
Discovery of Infrared Light in 1800
Herschel used a prism to
separate sunlight into colors.
He used a thermometer to
determine the temperature in
each color. (Two were placed
off to the side as controls).
The highest temperature was
found beyond red light (where
no light was seen).
Artists illustration from SIRTF web
page.
Historical motivation

Exploration & discovery


Technological opportunities


Neugebauer, Leighton, Low, Fazio, Townes
Bolometer (Low), PbS (Neugebauer), balloons (Fazio),
IR lasers & interferometry (Townes)
A few, key problems



Bolometric luminosities (Herschel, Johnson)
The Galactic Center (Becklin)
Star formation [many but especially Strom(s), Cohen,
Rieke(s) ]
IR Bolometer and Array Detectors
1.
Photon Detection in PN Junctions
- Review semiconductors
- The PN Junction
- Charge collection in PN junctions
Valence & Conduction Bands in Semiconductors
•When atoms (a) come together to
form a crystal, the outer energy
levels overlap and blend to create
bands (b).
•The outermost filled band is called
the valence band (c).
•Above the valence band, one finds
a forbidden energy gap -the “band
gap”, and (at higher energies)
conduction bands populated by
thermally excited electrons.
•In metals, the valence and
conduction bands overlap resulting
in conduction. In insulators, the
band gap is wider resulting in very
poor conduction.
Periodic Table
Semiconductors occupy column IV of the Periodic Table
(and have 4 valence electrons per atom)
P & N Type Semiconductors
•In a semiconductor, some electrons are
promoted from the valence band into
conduction by thermal excitation at room
temperature.
•These promoted electrons leave behind
positively charged “holes”.
•Both electrons in the conduction band,
and holes in the valence band, contribute
to conduction.
P & N type Semiconductors Continued

One can “dope” the semiconductor by adding impurities to the
crystal. Adding an impurity with more valence electrons than the
crystal will donate negative charges to the conduction band,
thereby creating an “n-type” semiconductor.

If the impurity has fewer valence electrons than the crystal, it will
donate holes to the valence band giving rise to a “p-type”
semiconductor.

When p-type material is butted against n-type material, the result
is a PN junction. In CCDs and most IR arrays that are in use
today, photo-excited charge is collected in PN junctions.
PN Junctions
•In a PN junction, positively charged holes diffuse into the n-type
material. Likewise, negatively charged electrons diffuse in the the ptype material.

•This process is halted by the resulting E field.
•The effected volume is known as a “depletion region”.
•The charge distribution in the depletion region is electrically
equivalent to a 2-plate capacitor.
Photon detection in PN junctions




A photon can interact with the semiconductor to create an electronhole pair.
The electron will be drawn to the most positively charged zone in the
PN junction, located in the depletion region in the n-type material.
Likewise, the positively charged hole will seek the most negatively
charged region.
Each photon thus removes one unit of charge from the capacitor. This
is how photons are detected in both CCDs and most IR arrays.
IR Arrays are “Hybrid” Sensors



A photosensitive array of PN
junctions is “bump bonded” to a
silicone readout multiplexer
(MUX).
This is done because silicon
technology is much more
advanced than any other
semiconductor electronics
technology. A modern MUX has
about as many transistors as the
most advanced Pentium CPU.
The “bump bonds” are made of
indium, a very soft metal used
for “welding” dissimilar
materials.
Schematic View of an IR Array



Note that each pixel has only one electrode.
Charge collection occurs in the depletion region near a PN Junction.
Charge is sensed in situ (it does not move as in a CCD).
Backgrounds
The Atmosphere
Astronomical Backgrounds
Atmospheric effects

Absorption



Emission



increased background noise
reduced integration times
Turbulence


reduced source flux
difficult calibrations
increased object size (“seeing”)
All effects vary with wavelength, time, altitude, line-ofsight
Atmospheric absorption versus airmass

The amount of absorbed radiation depends upon
the number of absorbers along the line of sight
AM=1
Atmosphere
I   I 0, 10
 mag / 2.5
AM=2
, mag  AM ,
where  is atm. extinction coefficient.
Atmospheric absorption versus 

Sharp cutoffs




defined primarily by
H2O
shape wavebands
Higher transmission
between lines with
higher resolution
Can introduce large
calibration errors for
low resolution
observations
(MNRAS, 1994, 266,
497)
Altitude: 4200m
Airmass: 1.0
H2O column: 1.2mm
Resolving power 3000
"These data, produced using the program
IRTRANS4, were obtained from the UKIRT
worldwide web pages.”
http://www.jach.hawaii.edu/JACpublic/UKIRT/astronomy/calib/atmos-index.html
Atmospheric absorption versus  - high
res
Array
defects
CO2 absorption
lines


R = / ~ 23,000
+
+
Keck II 10-m
Figer et al. 2000,
ApJ, accepted
Atmospheric absorption versus altitude

Particle number densities (n) for most absorbers fall off
rapidly with increasing altitude.
I  I0, e
 
,where   is optical depth,
    ndx   e


 x/ x0
dx
x0,H20 ~ 2 km, x0,CO2 ~ 7 km, x0,O3 ~ 1530 km
So, 95% of atmospheric water vapor is below the altitude
of Mauna Kea.
Atmospheric Transparency on Mauna
Kea
CSO web page
.
Atmospheric Transmission (0.9-2.6 mm)
Atmospheric absorption versus altitude
Telluric OH and Thermal Emission:
Mauna Kea
NIRSPEC
R~2000
J H K
Sky
Thermal
Background
OH Airglow: time variability
Atmospheric emission: Blackbody
Total power onto a detector:
P = h AW n esky Bn(Tsky)
h:
transmission of all optics x Q.E.
esky: emissivity of sky
A:
telescope area
W:
solid angle subtended by focal plane
aperture
n:
bandwidth
Bn(Tsky): Planck function
At 10 mm, typically: h ~ 0.2, e ~ 0.1, AW ~ 3x10-10 m2 Sr
n ~ 1.5 x 1013 Hz (10 mm filter), T ~ 270 K
P ~ 10-9 W or ~ 4 x 1010 g s-1
Atmospheric Turbulence

A diffraction-limited point spread function (PSF) has a
full-width at half-maximum (FWHM) of:
 FW HM 1.2

{m}
D{m}

D{m}
{" }
In reality, atmospheric turbulence smears the image:
 FW HM 0.25

{radians}  0.25
{mm}
{mm}
r0 {m}
{" }, where r0  6 / 5 .
At Mauna Kea, r0=0.2 m at 0.5 mm.
“Isoplanatic patch” is area on sky over which phase is
relatively constant.
Atmospheric Turbulence
1.4O seeing
0.5O seeing
no seeing!
Lick 3-m
Keck I 10-m
HST/NICMOS 2.4-m
Figer 1995
PhD Thesis
Serabyn, Shupe, & Figer
Nature 1998, 394, 448
Figer et al. 1999
ApJ. 525, 750
Background - sources
 Atmosphere


thermal
molecular
 Telescope


thermal
scattering
 Zodiacal
light
 Astronomical sources
Background - sources: Atmosphere
 Thermal
n
 1

Csky,thermal  hinstrhtele AW e sky Bn (Tsky )QE {e s }
hn
 OH


The average OH line intensity is approximately
25,000 g s-1 m-2 asec-2 mm-1.
The continuum between lines is about 50 times
lower than this value (in the H band).
Background - sources: Telescope scattering
 Mirrors
 s 2
Iscattered    , where s is RMS deviation
 
from a perfect surface.
 Baffle
edges and walls
 Secondary support
Background - sources: Astronomical
 Astronomical
objects can be objects of
interest or noise contributors, depending
on the project.
Sunlight,
moonlight
Light scattered by solar system dust
(“zodiacal”)
Light emitted (thermal) by solar system dust
(“zodiacal”)
Stars (especially in a crowded field)
Light emitted by interstellar dust (“cirrus”)
Background - sources: Astronomical
Radiation Processes
Absorption in Insulators: resonance features
Lattice resonances
np2 g n
e” = (n 2 - n2)2 + g2n2
0
log(e”)
0
Vibrational modes
~ 1 – 30 mm
-2
-4
s~
-6
2
1
8pa
Im(e'') ~ n2

0
log (n)
-1
kn ~ s(n) / mp
~ n2
-2
long wavelengths
Radiative heating: isolated particle
Particle radius, a (spherical; rapidly spinning)
Temperature, T
Distance, r
Absorbed radiative power: pa 2
L
4pr 2
Emitted radiative power: 4pa2 sT 4
Luminosity, L
L 1/4 -1/2
T= (
)
r
16ps
Using en for small particles: T ~ r -2/5 Pe  4pa2 n0  n-1 Bn(T) dn

cf L. Spitzer, Jr., Physical Processes in the Interstellar Medium, ch. 9.1
Thermal emission

spectral radiance, brightness, specific
intensity:
In = e cos  Bn(T) W m-2 Hz-1 sr-1
e  emissivity (dimensionless)
Planck (blackbody) function:
2hn3
1
Bn(T) =
c2 exp(hn/kT) - 1
Peak in nBn:
max(mm) =
Flux density from surface:
Total flux:
Fn = p Bn(T) W m-2 Hz-1
F = s T4
3674 K
T
W m-2
s = 5.67 x 10-8 W m-2 K-4
Planck Function

Assumptions



Uniform temperature source
Source is opaque
Mathematical description
2hc 2 / 5
B 
exp hc / kT   1
Emitting Area


erg
 Units :

o

2

s
cm
A
ster


h  Planck Constant  6.63  10 27 erg s
c  Speed of Light  3.00  1010 cm s-1
k  Boltzmann Constant  1.38  1016 erg K -1
cm
Temperatur e K 
  Wavelengt h of Light
T  Uniform
Computed Blackbody Spectra
2hc 2 / 5
B 
exp hc / kT   1
Rayleigh-Jeans
Tail
Wien
Law
Blackbody Curves
Wien Displacement Law
Blackbody peak wavelength inversely
proportional to temperature
 Find peak wavelength by solving:

dB
0
d

5 1 e
y
 y
where
2hc 2 / 5
B 
exp hc / kT   1
where
Numerical solution :
Wien Law:
 pkT  0.29 cm K
y
hc
pk kT
y  4.97
T  300 K   pk  10 mm
T  3000 K   pk  1 mm
T  30,000 K   pk  0.1 mm
Relative dust extinction
10
A / AJ
1
0.1
0.01
0.001
0.1
1
10
Wavelength (mm)
100
1000
IR Sub-mm Science
Current interest in infrared

High redshift objects



obs = 0 (1+z)
5000 Å  >1 mm for z > 1
Classical problems require infrared data
Obscuration by dust

A ~ -1.9  A2.2mm ~ 0.1 AV
S. Mathis 1990, ARAA, 28, 37.

Now important for:
 Galactic nuclei, esp. AGN (unified model)
 Starburst galaxies
 Young stars
J.
Current interest in infrared

Very low mass objects & extrasolar planets


Tplanet ~ 50 to 500 K
TBD ~ 900 – 2000 K
peak ~ 5 – 50 mm
peak ~ 1 – 5 mm
Extrasolar planets, brown dwarfs,
and circumstellar disks are optically
faint but infrared bright.
Structure of a protostar
after Stahler, Shu, and Taam 1980, Ap.J., 241, 637.
Young infrared star: W33 A
after Soifer et al. 1979, Ap.J.Lett., 232, L53.
NICMOS
Mass Loss from Evolved Stars - 1
 Broad





Scientific Goals & Key Objectives
Measure outflow characteristics for evolved stars
 Temperature, density, velocity, and composition
 Radial dependence for resolved sources
Understand molecular and dust chemistry in outflows
 Nonequilibrium gas chemistry
 Dust formation mechanisms and rates
Understand dynamical mechanisms driving outflows
 Radiative acceleration beyond a few stellar radii
 Adams & MacCormack (1935), Spitzer (1938)
Predictive model of mass loss from evolved stars
 Function of stellar age and initial stellar mass
 Feedback on interstellar structure and composition
Test stellar evolution models for evolved stars
 Nuclear reaction pathways
 Internal mixing mechanisms
Mass Loss from Evolved Stars - 2

Key Measurements




Molecular lines at infrared and millimeter wavelengths
 Over 50 species detected in IRC+10216
 Line ratios constrain temperature and density
 Line shifts and widths constrain velocity fields
 Isotopic abundance ratios constrain stellar models
Infrared dust features
 A few dust families (silicates, graphites, ices, etc.)
 Band strengths constrain dust chemistry
Angular resolution (10 mas)
 Resolves radial dependence of outflow characteristics
 Directly image clumps and general asymmetry
 Measure proper motion of clumps in nearest sources
Spectral energy distribution constrains unresolved sources
Mass Loss from Evolved Stars - 3
Sources

Samples

Resolved outflow sources
 Cursory literature search
• Supergiants (I, II) and Miras
• Stellar angular diameter >5 mas
IRC+10216
60 mas
R Dor
57 mas
W Hya
45 mas
a Ori
44 mas
20 <  < 40 mas
larger than photosphere 5 sources
23 sources
10 <  < 20 mas
 Also proto-planetary nebulae
49 sources
5 <  < 10 mas
Evolved stars in clusters
 Typical distance is 2 kpc
 Main sequence gives progenitor mass
 Interpret using detailed studies of resolved sources
 Outflows

Angular
Diameter
Mass Loss from Evolved Stars - 6
Tsuji, Ohnaka, Hinkle, & Ridgeway (1994, A&A, 289, 469)
Mass Loss from Evolved Stars - 7
IRAS 09425-6040
AFGL 4106
Molster et al. (2001, A&A, 366, 923)
Molster et al. (1999, A&A, 350, 163)
Mass Loss from Evolved Stars - 8
Cernicharo, Guelin, & Kahane (2000, A&AS, 142, 181)
Planetary spectra
2
4
H2SO4
CO2
0
2
Venus
Jupiter
2
H2O O3
Earth
0
0
Saturn
Mars
0
10
Relative, linear scales
20
30
0
Wavelength (mm)
10
20
30
Disks & infrared emission
102
RY Tau x 10
nFn (10-12 W m-2)
104
1
Vega
DL Tau x 2
102
10-2
9700 K
1
10-4
GM Aur / 20
10-2
0.1
1
10 100 1000
Wavelength (mm)
b Pic x 0.1
0.1
1
10 100 1000
Wavelength (mm)
Beckwith & Sargent 1996, Nature, 383, 139-144.
Circumstellar Dust
ASWG: Marcia Rieke
Vega Disk Detection

mm)
Flux* Contrast
(mJy) Star/Disk
11mm
2.4
1.5x107
22mm
400
2x104
33mm
1300
3x103
Reflected & emitted
light detected with a
simple coronograph.
*per Airy disk
NGST resolution at 24mm = 5 AU at Vega, > 10 pixels
across the inner hole
Waelkens et al. 1996, A&A, 315, L245.
200
Comet Hale-Bopp
6 Oct 1996
Fn(Jy)
100
Foresterite is a "primordial"
constituent of Solar dust
0
HD 100546
200
Fn(Jy)
Foresterite Mg2SiO4
100
PAH
0
10
20
30
Wavelength (mm)
40
HD 100546 - SWS and LWS : all components
6
4
Wavelength (µm)
PAH
(11.3 µm)
8
Hot & cold continuum
Total
0
Crystalline forsterite
Amorphous olivine
-50
10
Malfait et al. 1998, A&A, 332, L25
[ OII ] (157.7 µm)
2
[ OI ] (63.2 µm)
0
Pf g
Br a
Br d
5
H2O - ice (43.8 µm)
PAH
Crystalline pyroxene (40 µm)
H2O - ice
PAH
PAH (7.8 µm)
PAH (8.6 µm)
50
HD 100546
Stellar photosphere
Hot continuum
Cold continuum
Total
PAH
PAH (6.2 µm)
100
10
PAH (3.3-3.4-3.5 µm)
FLUX (Jy)
150
FLUX (Jy)
200
Short wavelength part - SWS
15
Pf d
250
FeO
Wavelength (µm)
100
Radio to IR Spectrum of Luminous IR
Galaxies
“K-correction”
increases flux
density for high-z
objects.
Carilli & Yun 2000,
ApJ, 530, 618
Mid-IR Observations of NGC1068
Imaging the
starburst
component.
(a) Mid-IR continuum.
(b) PAH emission. (c)
SCUBA 450 um on
PAH. (d) CO on PAH.
Le Floch et al 2001,
A&A, 367, 487
Mid-IR Observations of NGC1068. II
(Top) Decomposition of
Mid-IR spectrum into
AGN and starburst.
(Bottom) ratio of
unresolved flux to
extended (40”) and total
emission
Le Floch et al 2001,
A&A, 367, 487
Broad Band SED of 3C273
A large fraction of
the bolometric
luminosity is reemitted in the IRsubmm band.
Average spectrum of
3C273. Dashed line is
extended jet. Dotted
line is contribution
from host galaxy.
Turler etal 1999,
A&AS, 134, 89
Seeing through the dust in Cen A
COBE/DIRBE Image of the Sky
60 mm = blue; 100 mm = green; 240 mm = red. Hauser etal.
COBE/DIRBE Image of the Sky
Zodiacal light
removed. 60 mm =
blue; 100 mm = green; 240
mm = red.
Extragalactic
Background
(Galaxy removed).
240 mm image.
Hauser etal.
Cosmic UV to mm Extragalactic
Background
Cosmic background
can be produced by
warm M82-like star
forming galaxies.
Genzel & Cesarsky 2000, ARAA, 38, 761
NGST and the Future
Background - sources: NGST
NGST Backgrounds
1.E+04
zod. lgt.
1.E+03
e-/s/pixel
pm_scat
sm_scat
1.E+02
pm_therm
sm_therm
1.E+01
dark
readout
SPR_5
1.E+00
SPR_1000
1.E-01
1.E-02
0.1
1
10
Wavelength [ mm]
100
The Future: NGST
Near-infrared observing facilities
Facility
HST
SOFIA
IRTF
UKIRT
NGST
Keck
D
(m)
2.4
2.5
3.0
3.7
8.0
10.0
s
(")
D.L.
1.0
0.5
0.5
D.L.
0.5
Tm
(K)
290
230
273
273
40
273
e
Imager
0.05
0.15
0.12
0.10
0.03
0.10
NICMOS
FLITECAM
NSFCAM
IRCAM
TBD
NIRC
Spectrometer
(NICMOS)
FLITECAM
CSHELL
CGS4
TBD
NIRSPEC
Sensitivity of Future IR Facilities
5s Flux Limits in 104 seconds
1E-14
1E-15
SOFIA
1E-16
FIRST
nFn
 W/m 2)
1E-17
SIRTF
1E-18
ALMA
HST WFC3
1E-19
FAIR (8m)
1E-20
ST2010
1E-21
NGST
1E-22
0.1
1
10
Wavelength (m m)
100
1000
The End