High contrast spectroscopy
Michelson Summer School 2004
Bruce Woodgate, Goddard Space Flight Center
Topics
Spectra of known planets
Examples of current attempts for high contrast spectroscopy
How to do better
Integral field spectrographs
Why do spectroscopy?
• Spectroscopy is harder than pictures, so why bother? (The public
won’t get to see it! )
- For detecting an object, if we don’t introduce much detector noise or light
losses, we can get all the photons from an object at once, so it can be faster
than serial filter imaging.
- For distinguishing an object from background, such as a planet from the
scattered light of a parent star, or from background stars, spectroscopic
features can identify it. Need not wait years for proper motion.
- Characterize the object’s atmosphere: molecules and atoms, temperature,
density
- Radial velocity, orbit kinematics
Examples of planets spectra
Spectra of
planets (from
Karkoschka, Icarus,
1994, 111,174)
Neptune spectrum
(Melillo, 2000)
Earth extinction spectrum
(from Karkoschka, Icarus, 111,174)
Earthshine
spectrum
(Woolf et al, ApJ
574, 430)
Keck spectrum of molecular hydrogen around eclipsed star
HST/STIS slit position for brown dwarf spectrum
GL 229b
GL 229a (star)
GL 229b
(brown
dwarf)
Cool object spectra
(Schultz et al, 1998, ApJL 482,
L181)
GL 229b spectrum taken with
HST/STIS
Examples of current attempts for high contrast
spectroscopy
STIS “coronagraphic” format
Real instrument uses mirrors
CCD detector
Focal plane wedge
mask blocks star
with selectable width
Hubble space telescope
+ corrector mirrors
Pupil plane
mask covers
15% outer ring
Correction for spherical
aberration only
No mid-frequency wavefront
correction
No masking of secondary and
spider diffraction
Gas and Dust cloud condensing around AB Auriga
- Hubble/STIS coronagraphic image, PSF subtracted
Visible
X-rays
↓ HH 409C
HH 409A

←
new HH knots?→
STIS 1998 Coronagraphic Image of HD163296 (Grady et al. 2000)
placed next to Chandra 2003 X-Ray Image (20”x20”) of the same
field, courtesy of Doug Swartz (USRA/NSSTC), adjusted to the same
plate scale. As HH 409A and 409C have moved out of the field, we
believe the x-ray image has revealed one if not two new HH knots.
Jets and Older
Herbig Ae Stars
• The first indication that
the conventional
wisdom that nearZAMS intermediatemass stars don’t have
jets was wrong was
provided by HD
163296.
HD163296 STIS red and UV spectra
RED – G750L
H alpha S II
UV – G140M
Si III
Lyman alpha red-shifted
Lyman alpha blue-shifted
Spectroscopic formats onto 2-D detectors
Focal Plane
Feed to Spectrograph
Echelle
(two spectral dimensions)
Long Slit
(one spatial, one spectral
dimension)
Tunable narrow
band imager
(two spatial dimensions,
serial tune for spectral)
1
ImageSlicer
(two spatial, one spectral
dimension)
2
3
4
1
2
3
4
Detector
STIS G140M Long
Slit Observations
over two epochs
The
1999
image
is
discussed in Devine et al
(2000).
Lyman
α
observations
manifest a proper motion
asymmetry between the jet
and
counterjet,
while
observations in Si III
highlight an asymmetry in
ionization. Looking at Si III,
it is also evident that the jet
is decelerating as it moves
away from the star.
HH 409C
HH 409B
HH 409A
HD163296
H Alpha
Fabry-Perot in
coronagraphic
mode at
Apache Point
3.5-m
telescope
HD163296
(H Alpha – offband)
Fabry-Perot in
coronagraphic
mode at
Apache Point
3.5-m
telescope
HD100546 – HST/STIS visible coronagraphic image, combining 2 roll angles
•
Mapping the Environment of
HD
100546
Long slit spectrum
obtained in October
2001 along the
system major axis.
Visible
UV Lyman alpha – G140M
•
Mapping the Environment of
HD
100546
System minor axis
observed in June
2002.
HD 100546: the Inner 100 AU
Minor
axis
Major
axis
• Ly a, H2, and reflection nebulosity is enhanced
along the system minor axis to the NE of the star.
• Consistent with an origin in the envelope and not
in the disk.
The HD 104237 Bipolar Jet
•G140M long slit
spectroscopy of HD
104237 reveals a
microjet in Lyman a.
•The approaching jet
is along PA=342°,
with the counterjet
along PA=162°.
•The approaching jet
extends 1.05” from
the star, while the
counterjet can be
traced 2.6” from the
star.
Planet transiting star – HD 209458
HST/STIS photometry (from Brown et al, 2001, ApJ, 552, 699)
Sodium in
occulting
planet’s
atmosphere
HD 209458b
HST/STIS
G750M spectra
from
Charbonneau
et al, 2002, ApJ
568, 377
λ=1750A
HST/STIS
echelle spectrum
- format used for
long high
resolution spectra
- one long
spectrum is
chopped into many
short spectral
orders
λ=1150A
Supernova
1987A
red spectra
Wide slit spectra
(2 arcsec), to
show 2 spatial
dimensions and
one spectral.
One spatial
dimension is
convolved with
the spectral
dimension
How to go deeper
• Observe all the wavelengths, and all the relevant spatial points
simultaneously
• Maintain high efficiency and low background
Factors controlling signal to noise (1)
The signal to noise per resolution element (S/N) in a measurement is
(S/N) = Nsig / (Nsig + Nb)1/2
where
Nsig = number of counts from source
Nb = number of counts from background = Nscatt + Nsky + Ndarks +Nread
Nscatt = residual counts from star (after any suppression)
Nsky = counts from sky, airglow, our zodi, circumstellar, interstellar, etc
Ndarks = detector dark counts = Cdarksnpt where
Cdarks = detector dark rate
np = number of pixels per resolution element
t = total exposure time
Nread = equivalent counts from detector read noise = nrnp(nrms)2
For multiple reads,
nr = number of reads = t/t0
where
where
t0 = max exposure time to remove cosmic rays
(nrms)2 = equivalent detector read noise counts per pixel per read
Factors controlling signal to noise (2)
Converting to rates, N = Ct,
(S/N) = Csigt / (Csigt + Cscattt + Cskyt + Cdarksnpt + (t/t0)np(nrms)2 )1/2
(S/N) = Csigt1/2 / (Csig + Cscatt + Csky + Cdarksnp + (np/t0)(nrms)2 )1/2
Renaming, (S/N) = Csigt1/2 / Σnoiserate1/2
Total exposure time, t = (S/N)2 Σnoiserate / Csig
For an excellent coronagraphic spectrograph (TPF-C) observing a
faint planet, using a cold detector, Csig, Cscatt, Cdarks, are small.
Then Csky or Cread = (np/t0)(nrms)2 could be the dominant noise source
contributing to Σnoiserate.
Read noise is dominant with the best regular CCD (with nrms ~ 2
electrons), if (np/t0)(nrms)2 > FskyAeffΩδλ.
[Csky = FskyAeffΩδλ, Fsky (zodi)~2.5x108 ph/(cm2s.sr.μ), Aeff=4.3x104cm2, Ω=1.5x10-14 sr, np=9, t0=1000s]
For TPF-C parameters, read noise is dominant for δλ < 0.23μ (R~10).
So for spectral resolution elements narrower than broad band filters, zero read
noise photon counters are needed.
See also Lindler and Heap simulations later.
Spectroscopic strategies
High contrast spectroscopy
1) Reduce scattering source
- coronagraphy, nulling interferometry
2) Reduce other backgrounds (detector, sky)
- photon counting detector, high angular resolution
3) Observe source and background simultaneously
- subtract background from source under same conditions, eg seeing, thermal, pointing,
deformable mirror status
4) Select wavelength, resolution, polarization to improve signal to
background
- eg UV, narrow spectral band for gas line emission, polarization for scattered light,
IR for thermal emission
5) Include as many source photons as possible
- integral field spectroscopy
6) Observe background surroundings broadly to estimate background
at source position
- integral field spectroscopy, separate source and background with spectral
template
7) Observe reference point source under similar conditions
- PSF subtract
TPF Spectroscopic Requirements
Wavelength range
0.5 – 1.0 microns
Resolving Power
~70 (Nyquist sampled if read noise zero)
Spatial sampling
(i) Nyquist sample 6.0 meter diffraction limit at 0.5 microns (0.018 arcsec )
(ii) Elliptical: Nyquist sample 6.0 x 3.5 m diffraction limit at 0.5 microns
(0.018 x 0.031 arcsec )
Spatial coverage
Cover coronagraphic dark hole (= 1.8 arcsec square if have 96 x 96 DM
actuators)
TPF SPECTROSCOPY
- TRADE BETWEEN SLIT AND INTEGRAL FIELD SPECTROGRAPH
Property
Transmission (using prism)
Slit (towards star)
~0.8
IFU
~0.8
Roll alignment for point source Needed
Not Needed
Alignment (slit to star)
Difficult
Easy
Multiple planets?
No
Yes
Disk spectra
Less sensitive
More sensitive
Help find planets?
No
Yes, if buried in speckles
Dark hole edge
Traditional Integral Field Techniques
Focal Plane
Feed to Spectrograph
Detector
Lenslet
Array
Fiber
Bundle
Image
Slicer
1
2
3
4
1
2
3
4
Based on a figure from Content, 1998.
INTEGRAL FIELD UNIT OPTIONS
Lenslet array – concentrates and separates images to allow room to
interleave spectra
- examples: CFHT/Tiger, Oasis, SAURON, SNIFS, OSIRIS, MEIFU
- smallest and lightest
- high throughput
- large spatial format
- limited spectral elements, use for low spectral resolution or small spectral
range
2 Mirror image slicer array – rectangle to line reformat
- examples: MEIFU, MUSE, JWST/NIRSpec, SNAP, GNIRS, KMOS, SPIFFI
- intermediate weight and size
- high throughput
- small spatial format
- large spectral format
Fiber/lenslet array – rectangle to line reformat
- examples: SMIRFS-IFU, GMOS, SILFID, INTEGRAL
- large and heavy
- lowest throughput (fiber light losses)
- OK for very wide field ground-based MOS
OSIRIS
James Larkin (PI), Alfred Krabbe (Co-PI), Andreas Quirrenbach(PS),
Sean Adkins, Ted Aliado, Paola Amico, Matthew Barczys, George Brims, John Canfield, Thomas
Gasaway, Christof Iserlohe, Evan Kress, Ken Magnone, Nick Magnone, Michael McElwain,
Juleen Moon, Gunnar Skulason, Inseok Song, Michael Spencer, and Jason Weiss
•
Lenslet Array Integral Field Spectrograph
– Dissects arcsecond sized regions of the sky in 2
dimensions
•
•
•
•
Spectral resolution sufficient to take advantage
of low background between OH sky lines
(R=3900)
Full z, J, H, or K spectra with single exposure
(1700 pixels)
Very sensitive due to the suppression of
atmospheric emission lines, the lack of slit
losses and the low noise detector.
Size ~ 1.5 tons – >
– Vacuum chamber is ~ 1 m3
– About 200 kg is taken to 70 K
y
l
x
Design Summary
(from Larkin)
Lenslet
Array
Cold
Pupil
AO
Focus
Spectrograph
Collimator
Optics
Grating
Filters
R. I. Collimating
Singlet
R.I. Camera
Singlet
Pupil Plane
Focal Plane
Lenslet
Reimaging Optics
Camera
Optics
Detector
Lenslet
Array
Pupil Plane
AO Focus
(from Larkin)
1mm
• MEMs Optical’s design is fused silica, biconvex
elements. Thickness is 1.0 mm with EFL of 0.8
mm. Pitch is 250 microns.
– 72x72 lenslet square area centered in 1.5” diameter
circular substrate.
• 2-3 microns of rounding between elements (98%
fill factor)
• 2 micron alignment front to back
• Sub-micron accuracy of pitch
• 1% variation in EFL across array.
Example of microlens array
- for Supernova factory integral field spectrograph (SNIFS)
A symmetric magnifier for the spectrograph
Insertable convex mirror
Imager
mirror
Dark
hole
Beam from
coronagraph
Imager
field
Outer field angles
To spectrograph
Microlens
array
Microlens element of array for IFS
Microlens (eg: 250μ dia, f/4)
From
magnification
stage
Microlens-based Integral Field Spectrograph layout
Precede by
magnification stage
Microlens
(element of
array)
Insert for
modes 2
and 3
Insert for
mode 3
Focal plane
Mask –
insert for
mode 2
Change: Field lens now
precedes microlens
Insert for
mode 1
TPF prime planet/speckle spectral IFS data format
- at microlens focus, and projected onto detector, without disperser
FWHM of diffraction limit
Microlens
Spot at focus of
microlens
200 x 200 microlens
array.
20 pix separation.
5 spectra
interleaved, each
separated by 4 pix.
4k x 4k detector
TPF prime planet/speckle spectral IFS data format
- without disperser, showing detector pixel spacing
FWHM of diffraction limit
Microlens
Spot at focus of
microlens
Detector pixel
row spacing
200 x 200 microlens
array.
20 pix separation.
5 spectra
interleaved, each
separated by 4 pix.
4k x 4k detector
TPF prime planet/speckle spectral IFS data format
- with disperser
200 x 200 microlens
array.
20 pix separation.
5 spectra
interleaved, each
separated by 4 pix.
4k x 4k detector
TPF prime planet/speckle spectral IFS data format
- with prism disperser
TPF baseline, 100 pixels for R~70
200 x 200 microlens
array.
20 pix separation.
5 spectra
interleaved, each
separated by 4 pix.
4k x 4k detector
TPF prime and auxiliary candidate spectrographic capabilities
Suggested point design layouts for science discussion and prioritization. 0.5 -1.0 micron range. 4k x 4k detector. 6.5 m telescope.
TPF Baseline
Mode 1
Mode 2
Mode 3
Resolving power.
Pix/spectrum
~70
100
200 – 20000
100
3000
4000
3000
4000
Spatial Image points
Ang resolution (arcs)
Ang range (arcs)
200 x 200
0.037 x 0.037
3.7 x 3.7
200 x 200
0.037 x 0.037
3.7 x 3.7
200 x 5
0.037 x 0.037
3.7 x 0.185
33 x 33
0.22 x 0.22
3.7 x 3.7
Disperser
Prism
Grating
(tiltable). + Filter
Grating
Grating
Ensure other
grating orders
avoid detector
Microlens array
Concentration factor
# interleaved spectra
f/#s
200 x 200
20
5
500:4
200 x 200
20
5
500:4
200 x 5 (mask?)
20
5
500:4
33 x 33
120
30
500:4
Exchange
microlens array for
mode 4
Science capabilities
Planets/speckles
Circumstellar
disks.
Galaxy
dynamics
Galaxy jets
Stellar jets
SN in galaxy
QSO in galaxy
Galaxy pop
synthesis.
Gravity lenses?
High z galaxy
mergers
Stellar ejecta
(PN etc)
Comments
General astrophysics short high resolution spectral IFS data format
- with grating disperser and filter range blocker
Mode 1, 100 pixels for R~200 – 20,000, for selected short spectral regions
200 x 200 microlens
array.
20 pix separation.
5 spectra
interleaved, each
separated by 4 pix.
4k x 4k detector
General astrophysics long medium resolution spectral IFS data format
- with grating disperser
Mode 2, 4000 pixels for R~3000, for long spectra
Spatial format
200 x 5 microlens
array.
20 pix separation.
5 spectra
interleaved, each
separated by 4 pix.
4k x 4k detector
IFS microlens concentration factors
Concentration factor limitation contributors: diffraction, projection,
chromatic focus change, monochromatic aberrations.
Diffraction:
The FWHM at the focus of the microlens is given by dD=fλ/D, …. where D=
lens diameter, f=focal length.
Concentration factor, CD = D/d = D2/fλ …….for CD large, lenses should be large and fast.
(For OAO 250μ f/4, CD = 62, dD = 4µ)
Projection:
Concentration factor, CP = D/dp = D/fφ = θ/φ, ….where f/#in = φ, and f/#out = θ
(For a 10-m dia telescope (TPF), for a 250μ dia microlens to Nyquist sample the
diffraction limit diameter of 0.020 arcsec at λ = 0.5µ, need f/# = 500. Then CP = 500/4 = 125.)
Chromatic focus change:
Concentration factor CF= D/dc
For 250µ lens, from ray trace, best focus spot diameter = 2.4µ. Then CF = 100
Aberration:
From ray trace of spherical lens, best focus spot diameter = 4.0µ. Then CA = 62
Combined concentration factor: C = √(1/ΣCi2) = 38, for the example above (250μ f/4
lenslets with f/500 input). Then with 4 pixel spacing between spectra, can fit 9 spectra between image
points.
Alternative Reflective Image Slicer Integral Field Spectrograph
Detector
Slicer –
focussing
elements
Entrance
slits
Slicer – focal plane
Slicer Difficulties
• In a slicer spectrograph, one
spatial axis (along the slice)
is not sampled until the
detector, so all optical
components including the
grating introduce noncommon path errors.
• Each slice has different noncommon path errors.
• Introduced polarization is
also different for each slice.
From: Content, 98
Example Image Slicer detector format
- for 4 rows of re-imaging mirrors
Adjusting for the Elliptical Primary
- an asymmetric magnifier for the spectrograph
Asymmetric magnifier to sample 8 m x 3.5 m diffraction limit at 0.5 microns
Fixed convex cylindrical
mirror, horizontally
diverging
Beam from
coronagraph
Insertable convex
cylindrical mirror,
vertically diverging
Dark
hole
Fold
mirror
Outer field angles
Microlens
array
Imager
mirror
Dark hole
To spectrograph
Spectrograph field
Imager
field
Planet Detection Simulations
(from Lindler and Heap, GSFC, at Caltech TPF-C meeting June 2004)
• Four visits/two rolls per target
• Exposure times set to detect Earth-sized planet
at the scaled Mars distance. (limited to a
maximum of 14 days for all visits)
• Monte-Carlo simulation with random circular
orbits
 he = 0.1
• PSF reset every 10,000 seconds to reduce
speckle residual
• Plots color coded by stellar luminosity
• Spectra integrated over entire wavelength range
to compute S/N
6 x 3.5 Imaging Mode: Readnoise=3, Darkrate=0.001
6 x 3.5 IFU Rpower=10 Readnoise=3 Darkrate=0.001
6 x 3.5 IFU Rpower=10 Readnoise=1 Darkrate=0.001
Planet Detection via
Imaging vs. R=10 Spectroscopy
Imaging
• Standard CCD’s can
be used
• Requires Telescope
roll to remove
speckles
• Requires a very
stable PSF
Spectroscopy
• Requires detector
with readnoise < 1e• No telescope dithers
needed
• Planet confirmation
information obtained
• Reduced PSF stability
required