Lecture 14
Part 1: AO System Optimization
Part 2: How to be a savvy user
and consumer of AO
Claire Max
Astro 289, UC Santa Cruz
Feburary 21, 2013
Page 1
Optimization of AO systems
• If you are designing a new AO system:
– How many actuators?
– What kind of deformable mirror?
– What type of wavefront sensor?
– How fast a sampling rate and control bandwidth (peak
capacity)?
• If you are using an existing AO system:
– How long should you integrate on the wavefront sensor?
How fast should the control loop run?
– Is it better to use a bright guide star far away, or a
dimmer star close by?
– What wavelength should you use to observe?
Page 2
Issues for designer of astronomical AO
systems
• Performance goals:
– Sky coverage fraction, observing wavelength, degree
of image compensation needed for science program
• Parameters of the observatory:
– Turbulence characteristics (mean and variability),
telescope and instrument optical errors, availability of
laser guide stars
• AO parameters chosen in the design phase:
– Number of actuators, wavefront sensor type and
sample rate, servo bandwidth, laser characteristics
• AO parameters adjusted by user: integration time on
wavefront sensor, wavelength, guide star mag. & offset
Page 3
Example: Keck Observatory AO
“Blue Book”
• Made scientific case for Keck adaptive optics
system
• Laid out the technical tradeoffs
• Presented performance estimates for realistic
conditions
• First draft of design requirements
The basis for obtaining funding commitment
from the user community and observatory
Page 4
What is in the Keck AO Blue Book?
• Chapter titles:
1. Introduction
2. Scientific Rationale and Objectives
3. Characteristics of Sky, Atmosphere, and Telescope
4. Limitations and Expected Performance of Adaptive
Optics at Keck
5. Facility Design Requirements
• Appendices: Technical details and overall error budget
Page 5
Other telescope projects
have similar “Books”
• Keck Telescope (10 m):
– Had a “Blue Book” for the telescope concept itself
• Thirty Meter Telescope:
– Series of design documents: Detailed Science Case,
Science Based Requirements Document, Observatory
Requirements Document, Operations Requirements
Document, etc.
These documents are the kick-off point
for work on the “Preliminary Design”
Page 6
First, look at individual terms in error
budget one by one
• Error budget terms
– Fitting error
– WFS measurement error
– Anisoplanatism
– Temporal error
• Figures of merit
– Strehl ratio
– FWHM
– Encircled energy
– Strehl ratio
Page 7
Fitting error: dependence of Strehl on
 and DM degrees of freedom
Deformable mirror fitting error only
(
S = exp -s f2
)
5/3
é
ædö ù
= exp ê -0.28 ç ÷ ú
è r0 ø úû
êë
æ l ö
r0 ( l ) = r0 ( l = 0.5 m m ) ç
è 0.5 m m ÷ø
6/5
5/3
é
ö æ 0.5 m m ö 2 ù
æ
d
S = exp ê -0.28 ç
÷ø ú
çè
÷
l
è r0 ( l = 0.5 m m ) ø
úû
êë
• Assume very
bright natural
guide star
• No meas’t error
or anisoplanatism or bandwidth error
Strehl increases for smaller subapertures
and shorter observing wavelengths
Page 8
Strehl increases for smaller subapertures
and longer observing wavelengths
• Assume very
bright natural
guide star
Decreasing fitting error
• No meas’t error
or anisoplanatism or bandwidth error
Deformable mirror fitting error only
Page 9
Strehl increases for longer  and
better seeing (larger r0)
• Assume very
bright natural
guide star
Decreasing fitting
error
• No meas’t error
or anisoplanatism or bandwidth error
Deformable mirror fitting error only
Page 10
Wavefront sensor measurement error:
Strehl vs  and guide star magnitude
Assumes no DM fitting error or other error terms
s S2- H
æ 6.3 ö
ȍ
è SNR ÷ø
(
S = exp -s S2- H
2
)
é æ 6.3 ö 2 ù
= exp ê - ç
÷ø ú
è
SNR
êë
úû
SNR increases as flux from guide star increases
Strehl increases for brighter guide stars
But: SNR will decrease as you use more and more
subapertures, because each one will gather less
light
Page 11
Strehl increases for brighter guide stars
Assumes no DM fitting error or other error terms
bright star
Decreasing
measurement error
dim star
Page 12
Strehl vs  and guide star angular
separation (anisoplanatism)
é æ q ö 5/3 ù
2
ùû = exp ê - ç ÷ ú ,
S = exp éë -s iso
êë è q 0 ø úû
r0
q0 = µ l 6 /5
h
5/3
2ù
é æ
ö æ 0.5 mm ö
q
S = exp ê - ç
çè
÷ø ú
÷
l
êë è q 0 (0.5 mm ) ø
úû
Strehl increases for smaller angular offsets
and longer observing wavelengths
Page 13
Strehl increases for smaller angular
offsets and longer observing
wavelengths
0 arcsec
4 arcsec
10 arcsec
20 arcsec
Page 14
PSF with bright guide star: more degrees
of freedom ⇒ more energy in core
Peak intensity relative to diffraction limit
Point Spread Function very bright star, λ = 2.2  m, D / r0 = 8.5
1
218 DOF
50 DOF
0.1
24 DOF
12 DOF
2 DOF
(TIP-TILT)
uncorrected
0.01
Seeing
limited
0.001
0.0001
0
0.1
0.2
0.3
Radius (arcsec)
0.4
0.5
Page 15
What matters for spectroscopy is
“Encircled Energy”
Fraction of light encircled within diameter of xx arc sec
2.2 
Encircled Energy Fraction
Diffraction
limited
1.65 
1.25 μ
0.88 
uncorrected
0.7 
Page 16
Overall system optimization
• Concept of error budget
– Independent contributions to wavefront error from
many sources
• Minimize overall error with respect to a parameter
such as integration time or subaperture size
Page 17
Error model: mean square wavefront error is
sum of squares of component errors
• Mean square error in wavefront phase
s =s
2
p
2
Meas
Meas’t
+s
2
BW
+s
s
=s
+s +s
Fitting
Timelag
2
Meast
2
DM
2
S-H
2
iso
Isoplan.
é 3.5qb d ù
Ȑ
ë SNR l úû
2
T -T
+ .....
Tip-tilt
2
Page 18
Signal to Noise Ratio for a fast CCD
detector
Flux ´ Tint
Flux ´ Tint
SNR =
=
1/2
2
2
2
2
Noise
éës PhotonNoise + s SkyBkgnd + s DarkCurrent + s ReadNoise ùû
• Flux is the average photon flux (detected photons/sec)
• Tint is the integration time of the measurement,
• Sky background is due to OH lines and thermal emission
• Dark current is detector noise per sec even in absence of
light (usually due to thermal effects)
• Read noise is due to the on-chip amplifier that reads out
the charge after each exposure
Page 19
Short readout times needed for wavefront
sensor ⇒ read noise is usually dominant
• Read-noise dominated: read noise >> all other noise
sources
• In this case SNR is
SNRRN =
Flux ´ Tint
1/2
éë R n pix ùû
2
Flux ´ Tint
=
R n pix
where Tint is the integration time, npix is the number of
pixels in a subaperture, R is the read noise/px/frame
Page 20
Now, back to calculating measurement
error for Shack-Hartmann sensor
1
2p d ù
é p
2
@ ê
Jb
rad
l úû
ë 4 2 SNR
2
s
2
S-H
• Assume the WFS is read-noise limited. Then
SNRS-H
Flux ´ Tint
=
and
R n pix
dù
é 3.5q b d ù é
=ê
= ê 3.5q b ú
2
ú
l û ( Flux ´ Tint )
ë SNR l û ë
2
s
2
S-H
2
R 2 n pix
Page 21
Error model: mean square wavefront error is
sum of squares of component errors
s =s
2
p
2
Meast
+s
2
BW
+s
2
DM
R n pix
æ Tcontrol ö
dù
é
s = ê 3.5qb ú
2 +ç
l û ( Flux ´ Tint ) è t 0 ÷ø
ë
2
2
p
2
5/3
+ s + .....
2
iso
æ dö
+ mç ÷
è r0 ø
5/3
æqö
+ç ÷
è q0 ø
5/3
+ ....
Flux in a subaperture will increase with subap. area d2
Tcontrol is the closed-loop control timescale, typically ~ 10 times the
integration time Tint (control loop gain isn’t unity, so must sample
many times in order to converge)
Page 22
Integration time trades temporal error
against measurement error
æ 10Tint ö
dù
é
s » ê 3.5q b ú
2 +ç
l û ( Flux ´ Tint ) è t 0 ÷ø
ë
2
2
p
R 2 n pix
5/3
From Hardy,
Fig. 9.23
Measurement error
r0= 0.1 m
Temporal error
1 / t0 = 39 Hz
Optimum
integration time
Page 23
First exercise in optimization:
Choose optimum integration time
• Minimize the sum of read-noise and temporal errors by finding
optimal integration time
é
dù
2
R 2 n pix
s 2p = ê 3.5q b ú
l û ( Flux ´ Tint )2
ë
æ 10Tint ö
+ç
è t 0 ÷ø
5/3
2
ds 2p
R 2 n pix
dù
5 æ 10 ö
é
= 0 = -2 ê 3.5q b ú
+ ç ÷
2 3
dTint
l û ( Flux T int ) 3 è t 0 ø
ë
Tintopt
2
2
é
R
n pix ù
q
d
æ
ö
5/3
b
= ê 0.32t 0 ç
ú
÷ø
2
è
l
( Flux ) úû
êë
5/3
Tint2/3
3/11
• Sanity check: optimum Tint larger for long τ0, larger read noise
R, and lower photon Flux
Page 24
Similarly, subaperture size d trades
fitting error against measurement error
é
dù
2
R 2 n pix
s f2 = ê 3.5q b ú
l û ( Flux ´ Tint )2
ë
æ dö
+ mç ÷
è r0 ø
5/3
Flux = I ´ area = I ´ d 2 where intensity I = photons/(sec cm 2 )
2
R 2 n pix
æ dö
dù
é
2
s f = ê 3.5q b ú
+
m
çè r ÷ø
l û ( I ´ d 2 ´ Tint )2
ë
0
2
R 2 n pix
æ dö
1
é
ù
2
s f = ê 3.5q b ú
+
m
çè r ÷ø
l û ( I ´ Tint )2 d 2
ë
0
Hardy, Figure 9.25
5/3
5/3
• Smaller d: better fitting error,
worse measurement error
Page 25
Solve for optimum subaperture size d
é
dù
2
R 2 n pix
s f2 = ê 3.5q b ú
l û ( Flux ´ Tint )2
ë
æ dö
+ mç ÷
èr ø
5/3
0
Flux = I ´ area = I ´ d 2 where intensity I = detected photons/(sec cm 2 )
é
dù
2
R 2 n pix
s f2 = ê 3.5q b ú
l û ( I ´ d 2 ´ Tint )2
ë
æ dö
+ mç ÷
èr ø
0
5/3
2
R 2 n pix
1ù
é
= ê 3.5q b ú
l û ( I ´ Tint )2 d 2
ë
æ dö
+ mç ÷
èr ø
5/3
0
Set derivative of s f2 with respect to subaperture size d equal to zero:
dopt
ìï 6 r05/3 é 3.5q b ù 2 R 2 n pix üï
=í
êë l úû (I ´ T )2 ý
5
m
ïî
ïþ
int
3/11
dopt is larger if r0, read noise, and npix are
larger, and if Tint and I are smaller
Page 26
LGS (10th mag TT star) Case
Wavefront Error (nm)
Atmospheric Fitting Error
110
Bandwidth Error
146
High-order Measurement Error
150
LGS Focal Anisoplanatism Error
208
Uncorrectable Static Telescope Abs
66
Dyn WFS Zero-point Calib Error
80
Residual Na Layer Focus Change
36
High-Order Aliasing Error
37
Uncorrectable AO System Aberrations
30
Uncorrectable Instrum Aberrations
110
Angular Anisoplanatism Error
24
Tilt Measurement Error (one-axis)
11
Tilt Bandwidth Error (one-axis)
18
Residual Tel Pointing Jitter (1-axis)
96
Total Wavefront Error (nm) =
Strehl at K-band =
Assumptions
Zenith angle (deg)
Guide star magnitude
HO WFS Rate (Hz)
TT Rate (Hz)
Laser power (W)
WFS camera
Subaperture diameter (m)
Science Instrument
Amplitude of vibrations (arcsec)
d0 (m)
TT sensor
370
0.3
Keck 2 AO error
budget example
(bright TT star)
10
10
500
1000
16
CCD39
0.5625
NIRC2
0.2
2.41
STRAP APDs
Page 27
Summary: What can you optimize when?
• Once telescope is built on a particular site, you don’t
have control over 0, θ0 , r0
• But when you build your AO system, you CAN optimize
choice of subaperture size d , maximum AO system
speed, range of observing wavelengths, sky coverage,
etc.
• Even when you are observing with an existing AO
system, you can optimize:
– wavelength of observations (changes fitting error)
– integration time of wavefront sensor Tint
– tip-tilt bandwidth
– brightness and angular offset of guide star
Page 28
How to be a savvy user and consumer of
AO systems? Topics
• What kinds of astronomy are helped by AO?
• For users of astronomical AO:
– How to plan your observations
– What questions to ask when you get to the telescope
– Observing procedures
• For critical readers of AO papers in journals:
– How to assess the reality of AO results reported in the
literature
– Which data should you take seriously?
– What are “danger signs” that should make you doubtful?
Page 29
What kinds of observations will be
helped by AO? (1)
• See details that were not previously present
– Qualitative: can make new morphological statements
– Quantitative: need to know Point Spread Function; need to
understand PSF error bars
• Detect fainter objects/features
– Works for point sources
– But: IR AO systems inject more thermal background, because of
many mirror surfaces. Also throughput to detector goes down.
– In astronomy, faint extended objects can actually be harder to
see with AO. Limiting factor is sky background, and AO doesn’t
improve this for extended objects.
Page 30
What kinds of observations will be
helped by AO? (2)
• AO increases image contrast:
– Increased Strehl ratio ⇒ sharper edges, brighter features (if
they are close to diffraction limit)
– Detecting faint things close to bright things:
» companions to bright stars
» host galaxies of quasars
» stellar and protoplanetary disks
– Caution: Contrast improvement may not be helped by AO for
extended features, unless they have structure at λ / D
• AO permits more precise astrometry
– Can measure position of a point source more accurately if a) it
is smaller, and b) it is brighter
– But need other stars in the field to create a reference frame
Page 31
See new details and structure
Neptune, Keck, no AO
Neptune, Keck, AO
• Structure is dramatically clearer
• But can be hard to measure quantitative brightness of features
– AO PSF “spills” light from bright features into fainter areas
Page 32
Spilling of light, Neptune bright clouds
• Light from bright
compact cloud
region “spills”
over limb, and
into nearby dark
areas
• How do you tell
what the “real”
intensity is, in a
dark region?
Page 33
Will I detect fainter objects with AO? (1)
• Assume a point source under skybackground-limited conditions.
Total flux from object is Fobj
(ergs/cm2 or watts/m2).
• Generally choose size of pixel such
that two pixels sample a typical
point-source diameter. So within
the area of the PSF, npix ~4
• The area of the PSF on the sky is ~
(2λ/D)2 for AO, but is ~ (2λ/r0)2
without AO
• So if all else is the same, the sky
background Bsky within the PSF of a
point source is (D/r0)2 larger for the
no-AO case
SNRAO =
Strehl ´ FobjTAO
¢ t int
sky
n pix BAO
TAO
¢ t
(TAO
¢ = trans. (tel. + AO + instr.)
SNRsee =
FobjTNoAO
¢ tint
sky
n pix BNoAO
TNoAO
¢ t
(TNoAO
¢ = trans. (tel. + instr.)
2
æ Dö
sky
sky
BNoAO » ç ÷ ´ BAO
èr ø
0
Page 34
Will I detect fainter objects with AO? (2)
• Lick AO (  = 1.65 microns ):
S = 0.4
D=3m
r0 = 0.6 m
T’ao / T’noao = 0.5
• At 1.65 microns, the sky background per arc sec is the same with
and without AO, so
sky
AO
BAO
TAO
¢ n nopix
æ D ö TAO
SNRAO
¢
= Strehl ´
@
Strehl
@ 3.5 ´ Strehl
sky
AO
ç
÷
SNRseeing
Bno- AOTnoAO
¢ n pix
¢
è ro ø TnoAO
• If Strehl is < 0.3, AO doesn’t give sensitivity advantage even for
point sources
Page 35
Galactic Center: NGS to LGS AO
Credit:
Andrea
Ghez’s group
at UCLA
Best NGS
Page 36
Galactic Center: NGS to LGS AO
LGS
Page 37
Structure in extended objects
• If goal is to see diffraction-limited
structure, we need to achieve good
signal to noise in each pixel
SNRAO =
• Overall transmission T’AO ∝ D2
• AO pixels are usually diffractionlimited, so npix=aobj/apix= aobj/ ( /D)2
SNRAO µ
• Can increase SNR by binning pixels
– If object is diffuse (aobj >>  /D)
don’t need diffraction limited
pixel sampling anyway
n pix BtotalTAO
¢ t int
Strehl ´ Fobj D t int
aobj
( l / D )2
• SNR with AO indep. of telescope
diameter!
• The larger the object (aobj) the
lower the SNR per pixel
Strehl ´ FobjTAO
¢ t int
SNRAO
æ Strehl ´ D 0 t ö
int
µç
÷
aobj
è
ø
Page 38
AO yields higher contrast, for small
features
• T. Rimmele
• AO for imaging
surface of Sun
AO off
AO off
• Higher contrast
on bright
granules, dark
regions in
between where
B field is
emerging from
sub-surface
Page 39
Example of higher contrast: vision
science images of human retina
• Austin Roorda and David Williams
Without AO
With AO: resolve
individual cones
Page 40
AO yields better contrast for faint
objects next to bright objects
Page 41
AO can permit more accurate astrometry
(precise position measurement)
• For a point source, accuracy of centroid measurement
increases with intensity, decreases with image size
• AO helps both of these:
AO
• But: need stars with
known positions in field
of view. AO field of view
tends to be small
No AO
Binary stars are
perfect for relative
astrometry
Page 42
Questions? Discussion?
Page 43
How to plan observations ahead of time
• First requirement: understand what Strehl ratio you
will need for your science project to succeed
• Estimate exposure time needed to achieve good SNR
– Some AO systems have exposure time calculators
– Or check with folks who have observed in the past
• Refer to web pages to see what brightness guide star, at
what distance, at what zenith angle, you will need
• Check AO system web page for maximum offset
between science target and guide star
• Search star catalogues to find guide stars
Page 44
Star catalogs for guide star search
(After B. MacIntosh)
Catalog
Mag
Limit
Spectral
info?
Notes
SAO/PPM
9
* Types
Old catalogue, bright stars
Good for quick PSFs
Hipparcos
9
Colors
Very accurate, catalogue available
as IDL file
1112
Colors
Very accurate
HST Guide
Star
15
None
Unreliable (but good near bright
stars), locally searchable in IRAF
USNO
B1.0
20
Colors
Incomplete near bright stars,
funky close to big galaxies
Tycho 2
Page 45
Finding a guide star: Tools
• VizieR
http://vizier.u-strasbg.fr/viz-bin/VizieR
– Has the ability to do constrained searches – limited in
position, magnitude, etc. – from a list of input targets
– Results can be read into IDL or spreadsheet for sorting and
processing
• Aladin (one of Vizier’s capabilities)
– http://aladin.u-strasbg.fr/java/nph-aladin.pl
– Can overplot a Digital Sky Survey image of your target with
all the stars it can find from the USNO B1 catalogue. Very
useful for finding guide stars.
– Gives B and R magnitudes of all USNO stars.
Page 46
• Aladin and
USNO B1
catalog: virtues
and pitfalls
• Great user
interface, many
surveys
• But gets
confused near
galaxies,
nebulosity
• Check out
potential guide
stars by eye!
Page 47
Some
observatories
have their
own online
guide star
tools
Page 48
Other questions to address prior to
observing with AO system
• PSF calibrations
– What is my PSF star calibration strategy?
– What is the impact of anisoplanatism?
• Observing time
– Calculate exposure times needed for good SNR
– Have I accurately estimated AO’s overhead (wasted time)?
• Calibration and flat-fielding issues
– How will I calibrate the sky fluxes (offset skies, dithering, other?)
– How will I calibrate detector response variations?
– How will I calibrate photometry (brightness measurement)?
» Usually observe photometric standard stars
» How often? In what sequence?
Page 49
“PSF stars”
• Before, after, and sometimes intermingled with
observing science target, observe “PSF star”
• Constraints:
– If science target is offset in angle from guide star,
can find PSF star pair with similar relative offset
– Should be at ~ same zenith angle as science target
(but typically an hour or two earlier or later)
– PSF star should produce same number of wavefront
sensor counts as guide star for your science target.
• In practice it’s hard to meet all these conditions
• With LGS, I typically end up using the tip-tilt star as PSF
Page 50
Page 51
Sometimes you find creative endeavors
on the web (!)
Page 52
Laser guide star observing requires more
preparation
• US observatories have to submit target list to US Space
Command (satellite avoidance) in advance
– Not good form to destroy the detector on a billion
dollar satellite
• Specific formats required
• Check web pages for instructions
Page 53
Page 54
Questions to ask when you get to the
telescope
• If possible, come a day early and watch the previous
night’s observers use the AO system
• Ask for a “lesson” in how to control the AO system
from the instrument interface
• Typically the AO system is calibrated each afternoon
– Observatory staff will use an internal light source to
measure non-common-path errors every day (before
observing)
– Instructive to watch this process if you’ve never
seen it before
Page 55
Why we must calibrate for non-commonpath errors
To near-IR camera
Schematic of Lick AO system (one generation ago)
Page 56
Overview of the calibration process
(usually done by observatory staff)
• Close dome, lights out, flatten the deformable mirror
• Turn on internal light source (e.g. optical fiber with
diode or laser light)
– Record centroid positions on wavefront sensor
– Record image of internal reference on camera
• Adjust deformable mirror shape until image of
internal reference has highest Strehl ratio
• Record new positions of centroids on wavefront
sensor. These will be the “reference centroids” to
which AO loop will control.
Page 57
AO tuning for your guide star
• Wavefront sensor camera frame-rate and AO control loop gain
optimized for your guide star
– For fainter guide stars, want slower frame rate
– Typically need 100-200 counts per subaperture per wavefront
sensor frame, for good AO performance
– For fainter stars, use lower control loop gain (lower bandwidth)
• AO operator will take a sky background measurement for the
wavefront sensor
– Subtracted from each wavefront sensor frame
• Based on number of wavefront sensor counts, AO operator will run
a program to optimize the AO system performance (trade frame
rate against counts on wavefront sensor)
• Then offset from PSF star’s guide star to the PSF star itself, turn
on AO system, take images or spectra
Page 58
Re-tuning the AO system during the night
• When does operator re-tune AO system?
–
–
–
–
At each new telescope pointing
When background changes (clouds, moon)
When flexure changes (after slew, long integrations)
Whenever observer requests an updated tune-up
• I usually keep an eye on the number of
wavefront sensor counts per subaperture
– When it drops considerably below its original value,
ask for a re-tuning
Page 59
Other observing procedures are same as
for any infra-red observations
• Take sky backgrounds
– Necessary in IR: science target can be dimmer than
the sky background!
– Can nod to sky so that your science target is entirely
off the detector, or
– If your science target is small enough, can get sky
bkgnd just from dithering target on detector
• Observe photometric standard stars several times
during the night (if needed)
Page 60
Other observing procedures are same as
for any infra-red observations
• Dithering and nodding:
1
2
• IR array non-uniformity
requires sky measurement and
subtraction
• To obtain a sky subtraction,
usually need a multiposition
dither (1-2-3-4 etc.)
– If your science target is big,
good to get a separate sky
too
3
4
5 (sky)
Page 61
Questions? Discussion?
Page 62
How to assess the reality of AO results
reported in the literature
• Which data should you take seriously?
• What are “danger signs” that should make you doubtful?
Page 63
Taking data seriously: Three big issues
1. Strehl ratio and variability
2. Effect of using a non-point-source as a guide star or
tip-tilt star
3. Signal to noise ratio
Page 64
Taking data seriously: Three big issues
1) Strehl:
– Don’t trust lowStrehl results
Higher Strehl ratios
are more stable
– How low is low? My
rule of thumb:
“low” is S < 10%
– Problems: unstable
photometry,
variable PSFs
Credit: J. Christou et al.
Page 65
Big issues, continued
2) Finite-size object used as “guide star”
– Frequently produces artifacts on point spread
function
– Sometimes “double-star” PSF
– Look for independent measurement of PSF if possible
– Also there can be issues with using finite-sized
object as tip-tilt star
» Most important example: using bright nucleus of a
galaxy as the tip-tilt reference
» The more point-like it is, the better
» No firm rules here about what to do – try it!
Page 66
Big issues, continued
3) Signal to noise ratio of AO image or spectrum
– Rules of thumb (Hardy):
» SNR needed to recognize an object in a noisy
background: SNR > 5
» SNR needed for spectroscopy is much larger: people
use numbers like 20, 50, 100 per resolution element
(depends on the application)
Be sure to look carefully at section of published paper
where SNR is discussed.
– If it isn’t discussed, try to estimate it yourself.
Page 67
“Journal of Irreproducible Results”
• Danger signs:
• Low Strehl ratios (e.g. 5% - 15%)
• Use of an extended source as a “natural guide star”
– Can give PSFs that are double, or that have several
lumps
• Use of a “guide star” that IS a point source, but that is
embedded in a fuzzy region
– Also can give odd PSFs
• Look for repeatable independent measurements of PSF
Page 68
Radio galaxy 3C294 seen with UH AO
system
Diffraction spike
from guide star
(a double star?)
Stockton et al.
UH AO System
CFHT Telescope
Page 69
3C294 images from Hubble, Keck AO
Hubble (0.7 micron)
Keck AO (1.6 microns):
Bright point-like core,
plus fuzz on right
From Wim de Vries. LLNL
Page 70
Example of dangers from extended guide
star: Frosty Leo nebula
• UH AO system
• Closed AO loop on
one of the big blue
blobs
• Concluded central
star is double
• Not confirmed by
subsequent
observations
Page 71
ESO’s VLT MACAO Observations of
Frosty Leo
• Is it a
binary star?
Page 72
Conclusions
• AO systems can yield flakey results if:
– Guide star is extended, or too faint
– Strehl is too low or too variable
• Need good signal to noise (but that is no different from
“regular” observations)
• Need thoughtful preparation before an observing run
• But…. RESULT CAN BE WORTH THE TROUBLE!
Page 73