A small space mission designed to carry out a widefield UV transient survey
System Engineering - Report
1
Project Summary
2
Why go for mini-satellites?
 Why go for small satellites? NASA/ESA launch large (tons),
expensive (billions of $), slow (decade construction) missions
(e.g. JWST, EUCLID). Our goal is to do competitive science
with an agile program of light (few 100kg) and cheap (few
$10M) satellites.
 This is possible since:
 Technological advances provide powerful capabilities with modest
mass;
 Israel is a leader in this area, use IAI universal bus heritage;
 Recent years have seen increased "space access", with new
players (e.g. China, India, commercial) providing much lower cost
launch & communication.
3
(How) Can we beat larger missions?
 Large satellites: high resolution, high sensitivity, very small fields of view
(sub-degree). Our idea:
 Compromise on resolution & sensitivity in order to construct a small satellite
with a wide field of view (thousands of squared degrees)
 Identify rare transient events, which large satellites miss, follow up and
distribute (in real time) the location of the transient to larger space & groundbased observatories.
 Why UV?
 The transient UV/X-ray sky has not been explored and holds great prospects
for scientific discoveries.
 The technology for building light-weight wide-field (“Lobster Eye”) X-ray
optics is not mature enough. We have therefore decided to examine the
possibility of a wide-field UV mini-satellite.
 Prospects. If we are successful, the current mission may open the way
to, and be the first of, an agile program of small satellites doing
competitive science.
4
Wide field UV: Requirements
 Superseding earlier experiments.
 Our sensitivity goal is ~10 times less than that of GALEX (SNR=5 for ~0.01
photons/cm2s within Dl=0.044m at l<0.35m and 300s integration).
 Our field-of-view (FOV, ΔΩ) goal is, on the other hand, ~1000 times larger
(>1000 square degrees compared to 1).
 If the above requirements are met, the detection rate would be 30 times that
of GALEX.
 Guaranteed events.
 Since a wide-field UV survey has not done before, we expect unexpected
discoveries, or at least detecting some so far undetected types of events
(e.g. stellar disruptions by Black Holes, NS2 mergers).
 However, there are also some "guaranteed" transients: supernova shock
breakout events, which would be detected and would provide important
science output at low risk. The FOV and sensitivity requirements listed above
were chosen to provide more than 10 detections of supernova shock
breakouts per year.
5
Technical feasibility.
 Our study so far suggests that the requirements may be met
using existing technology. This is based on two main
arguments:
 The sensitivity is 10 times worse than that of the old GALEX;
 The technical estimates summarized in the rest of this
presentation imply that the goals may be met with
reasonable size telescope and detector.
 The Israeli (IAI) universal bus capabilities supersede the
Weight/Power/Comm./Stability requirements
A reduced capabilities/cost version may be chosen.
6
SN Breakouts: I. Scientific Background
 The explosion mechanism of SNe is not fully understood
 A major goal: Identify progenitor properties
 SNe usually detected days to weeks after explosion
 Detecting the “shock breakout” from the stellar edge
provides unique new constraints (eg progenitor radius,
envelope composition)
 Breakout: X-rays for 10’s of sec, UV for hours/day
 A handful of events have been observed at this early
stage
7
SN Breakouts: II. Flux, implied sensitivity
 UV/O (post) breakout emission
[form Rabinak & Waxman 11]:
-0.45
H envelope : TC.  1.9 R131/ 4tday
eV, Lbol.  10 43
E 51
-
R13tday
erg/s,
M / M Sun 2 / 3
  0.35(0.16) for radiative(convective)
C/O :
-0.35
TC.  1.8 R120.2tday
eV, Lbol.  5 10 42
 f UV (10hr, Dl/l  0.2) ~ 0.1
E 510.7
M / M Sun 
2/3
-0.1
R12tday
erg/s
R12
photons/cm 2s
2
D30 Mpc
 UV background in the 0.1-0.2μ range is ~2x10-8 erg/cm2/s/Å/sr
1 / 2
D
 q AT 
 f UV,5  0.03q  41 2 
photons/cm 2s
0.5'
 10 cm s 
A=area, T=integration time,
q=10-1q-1 is the overall (quantum + filters) efficiency of the detector,
qPSF-1 is the fraction of the flux that falls within the pixel
1
psf
8
Proposed System – Summary (1)
 Eight identical telescopes, each with:
 Aperture diameter
120 mm
 Focal length
290 mm
 F/#
F/2.4
 Field of view
12.1° x 12.1°
 IFOV
21.3 arcsec

Plate scale
 Percent energy/pixel
 Spectral band
 Filter type
 Visible suppression
710”/mm
>75% (result of PSF)
220-270 nm
Reflective (two in series)
2 x 10-3 from 300 to 1100 nm
9
Proposed System – Summary (2)
 Detector type
 Pixel array
 Pixel size
 Array size
 Binning
 Effective size
CCD
4096 X 4096
15 x 15 µm
61.4 mm square
2x2
30 x 30 µm, 2048 X 2048 pixels
 One binned pixel = 21.3 x 21.3 arcsec
 Quantum efficiency
60% average over band
10
Proposed System – Summary (3)
Performance
 Total field of view
 Fraction of sky covered
 Detection threshold
 Limiting AB magnitude
 Diffuse background
8 x 146 = 1170 sq. deg.
2.8%
0.006 ph/cm²/sec
18 at 300 secs integration time
0.08 ph/cm²/sec/(‘)² (assumed)
11
Proposed System – Summary (4)
Orbit
 Sun synchronous polar orbit
 LTDN
06:00 hrs
 Altitude (depends on launch possibilities)
 Minimum
 Desirable
 Inclination
 Stability
720 km
>1400 km
depends on altitude
(e.g. 8° for 720km)
Better than 50µrad in 300 sec
12
Status Summary
 A design meeting system requirements was reached





(SNR=5 for 0.006 photons/cm2s at 300s integration,
FOV>1000 squared degrees).
Constraints: Detector size (61 mm) and F/# (≥2.4)
Field of 12° x 12° gives better performance than 20°x 20°
originally proposed, because of larger lens diameter, despite
smaller field.
8 telescopes doubles detection rate, still within limit for no
direct view of Earth. However need baffles to prevent stray
light.
Reflective filter appears to offer acceptable sensitivity.
Higher orbit (>1400 km) has advantages of avoiding eclipse
and shorter baffles, but communication limitations.
13
Design Considerations
14
Challenges
 Field of view (FOV) which sensor can observe continuously is only a





small fraction of celestial globe
Signals are very weak and detection requires high sensitivity and
long integration times
Sensitivity limited by collection area of optics, by diffuse sky
background and detector sensitivity and noise, etc.
Selection of orbit is a compromise between best performance and
low cost. Orbit must allow virtually uninterrupted observation
Some periods of eclipse (loss of power) inevitable unless orbit is
above 1400 km
Stray light from the Earth could increase background, lower
sensitivity, in parts of orbit at different times of year. The higher the
orbit, the easier to minimize problem: needed baffles can be shorter
15
Challenges -continued
 Existing UV space sensors have small FOV which can avoid bright
stars or dense regions. Wide FOV means high photon rates from
stars and background. Classical image-intensifier detectors cannot
handle such high rates
 During part of the year, the Milky Way will cover the field of view and
for at least part of this time, the system will be inoperative due to very
high background or complete obscuration
 Distinguishing transient events requires comparison of image with a
reference image taken earlier. Signal processing is needed to
accomplish this
 Communication limitations probably mean such processing must be
done on-board
16
Telescopes and Field of View
 FOV is given by





FOV  tan 1 ( wdet / fl ) where wdet is width of
detector and fl is focal length of telescope
Collecting area is A  Dm2 / 4  fl 2 / F where Dm is diameter of
optical aperture and F is relative aperture (F/#)
F/# less than 2.4 is not practical in this system
Detector width of 60mm is best available with right
characteristics
With these, FOV of 40° x 40° would give A = only 10cm²
By dividing field into a number of telescopes, each with 12° x
12° FOV, we get A = 113 cm² for each
17
Coverage
 One telescope with 12.1° x 12.1° FOV = 146 sq.degrees
 Our choice: 8 telescopes – 1152 sq. degrees = 2.8% of the sky
 If we can detect an event at a level of 0.0056 ph/sec/cm², we
can expect to detect 240 * 0.028 = 6.7 SNe/year
 We believe this sensitivity can be achieved
18
Parameter relations
Supernova Breakout Detection rate
Detection _ rate  N t * N
Nt
Npixels
d
F
Dm
qPSF
q
T
Dl
Ibgnd
2
pixels
2
 qPSF
* q * T * Dl 
d

Dm

F 
I bgnd

3
4
- Number of telescopes
- Number of pixels
- Pixel size
- Optics F number (f/Dm)
- Entrance pupil diameter
- Percentage of energy on a pixel
- Detector quantum efficiency multiplied by overall transmittance
- Integration time
- Spectral region bandwidth
- The diffuse UV background flux
The performance based on this relation is shown in slide # 54
19
Design Study
20
1. Detectors
21
Detectors
 Image intensifier detectors with semi-transparent photo-
emissive cathode (such as GALEX) have:
 UV sensitive only but low Quantum Efficiency (typically 8%)
 Can only handle low photon rates (few 1000/s to 100,000/s)
 Generally round photocathodes up to ~ 60mm diam
 Spatial resolution limited
 No dark current
 Silicon CCD detectors have:
 Up to 60% QE, but
 Sensitive to visible also – need filter to suppress this
 Can handle millions of photons/sec
 4k x 4k arrays of 15µm pixels – 61 mm square
 Dark current and readout noise need to be reduced
22
Detectors - continued
 To keep dark current low, need cooling to ~230°K
 e2v IMO (Inverted mode operation) detector better than Non-IMO
 Use 2 x 2 binning (30 x 30 µm) to match optics, still resolution better
than needed (20 arcsec)
1.E+05
1.E+04
1.E+03
electrons/sec/pixel
1.E+02
1.E+01
1.E+00
1.E-01
NIMO
1.E-02
NIMO Binned
1.E-03
IMO
1.E-04
IMO Binned
1.E-05
1.E-06
1.E-07
1.E-08
170
190
210
230
250
270
290
Temperature [°K]
 A few bright stars will saturate pixel; charge will spread to a few
surrounding pixels
23
Spectral Response
 Preliminary e2v data indicates that 60% QE at 240nm is
possible, but visible response is high
 QE of e2v CCD231-84 array (provisional curve)
QE // -50°C // Si: 16µm // Astronomy Process // Process Temperature
Dependence ON // HafOx(1) = 18nm
100%
90%
80%
70%
60%
50%
40%
30%
20%
10%
0%
200
250
300
350
400
QE
450
500
550
600
Reflection
24
Visible response suppression
 In telescopes to map UV stars, like GALEX (or TAUVEX!)
response to visible must be much lower than to UV (because
visible spectrum much more intense than UV)
 In LIM transient sensor, some visible response tolerable as it
only adds somewhat to bright star signals which must be
ignored anyway
 To limit addition to background noise, out-of-band response
should be <10-3 of in-band response to spectrum like sun.
25
2. Filters
26
Filter manufacturers Materion, Acton
 Filter proposed by Materion ( formerly Barr) has >30% transmittance
in UV, visible transmittance is 10-4 for solar spectrum photon flux
Transmittance
Optical Density
0.5
7
6
0.4
5
0.3
OD
4
0.2
3
2
0.1
1
0
0
200
250
300
Wavelength [nm]
350
200
400
600
800
1000
1200
Wavelength [nm]
 Acton proposed a standard filter (see next slide)
27
Visible blocking filter
 Comparison of transmissive filters proposed so far
60
% Transmittance
50
JDSU 2
40
Acton
Materion
30
JDSU 1
20
10
0
190
210
230
250
Wavelength [nm]
270
290
 Materion seemed to be the best (before reflective filter proposed)
28
Filter manufacturer JDSU
 JDSU was paid to carry out a design study after indicating that




they could achieve high transmission
Their first proposal (JDSU1) was totally unacceptable, due to
misunderstanding of blocking needed
Second proposal, JDSU2, was better but still less than 30%
effective transmittance
However, they say that two reflective filters in series could
offer 95% transmittance. The filters would have to be at 45° to
optical axis to fit in system
At this angle, some reflection of polarized light in blocking
region but this can be tolerated (see slides 30, 40)
29
Reflective filter at 45° (JDSU Data)
100
90
80
% Reflectance
70
60
OAR
50
Overall Reflection S pol
Overall Reflection P pol
40
30
20
10
0
200
300
400
500
600
700
Axis Title
800
900
1000
1100
30
Reflective Filter
 Transmission of in-band UV is >95%, compared to 30% for best
transmissive filter. This would increase SNR by a factor of 1.8
 Reflectance of S-pol is ~ 14%. Two in series would be ~2% for half of
polarized light which would give very high “red leak” (~1%) in visible.
Diffuse background would increase by 150%, meaning 40%
reduction in SNR. Overall gain factor of only 1.12
 If the two are crossed so S-pol in first becomes p-pol in second, total
leak should be less than 0.2% which is acceptable. Increase in
diffuse background should not be more than 25%, meaning reduction
in SNR of 10% or less. Overall gain factor thus 1.6
 Reflective filter has also wider bandwidth (60nm). This may increase
performance although higher background, wider PSF could reduce
the gain from this factor (remains to be analyzed).
31
3. Optics
32
Optics
 Designing an wide field, low F/# optical system for UV is very





challenging
Catoptric (reflective) systems for wide field are complex
(dimensions, alignment, stray light)
For dioptric (refractive) systems for space, very few suitable
materials to enable correction of chromatic aberration. Wide
(40 nm) bandwidth adds to problem
Low F/# presents major challenge
PSF and percent energy on pixel it implies is a crucial
parameter
Study was made to compare options:
33
Optics trade-off study
Selected parameters:
 Spectral band 220 – 260 nm
 Effective focal length
290 mm
 FOV with 61.4mm detector 12.1° x 12.1° (diagonal ± 8.5°)
 Entrance pupil diameter 120 mm
 Effective resolution 20arcsec (30µm)
34
Optics trade-off study
 Catoptric objective – Three mirror anastigmat (TMA)
Optical layout
300 mm
Primary
Mirror
Perspective view
Secondary
Mirror
440 mm
Tertiary
Mirror
Image Plane,
(tilted w.r.t
central Chief-Ray)
MM
TMA F/
FFOV=
X
TMAf
c
A.N
MM
TMA F/
FFOV=
X
TMAf
c
Scale
A.N
35
Catoptric option – cont’d
 To prevent only direct stray light, long baffle required. (For full stray
light prevention, would need to be even longer)
17.2
Baffles
13
770 mm
Clear Aperture 320 mm
MM
TMA F/
FFOV=
X
TMAf
c
A.N
 Multiple telescopes would require huge assembly
36
Catoptric option – cont’d
TMA F/
TMAf
 Performance
FFOV=
c
X
Y
X
Y
X
Y
X
Y
X
Y
X
DIFFRACTION MTF
A.N
DIFFRACTION LIMIT
)
DEG
)
DEG
(-6.00,0.000)DEG
WAVELENGTH
NM
WEIGHT
(-6.00,6.000)DEG
)
DEG
DEFOCUSING
M
O
D
U
L
A
T
I
O
N
X
Y
Minimum NOMINAL MTF in the
entire FOV at Nyquist frequency
is 50%
Nyquist frequency for
pixel size of 18mm
SPATIAL FREQUENCY (CYCLES/MM
 PSF - % Energy on 30µm pixel very good
Field angle [°]
XAN
YAN
0
0
% Energy
On pixel
97.0%
0
6
98.8%
0
6
-6
-6
81.5%
81.8%
37
Dioptric Objective
 Preliminary
optical layout
MM
EFL=
D=
f=
LIM
S
Scale
A.N
 PSF - %energy on pixel
70% - 80%
% of energy on a pixel
Percentage of energy on a pixel of 30mmX30mm
90%
85%
80%
75%
70%
65%
60%
55%
50%
45%
40%
35%
30%
25%
20%
15%
10%
5%
0%
0.0
0.5
1.0
1.5
2.0
2.5
3.0
3.5
4.0
4.5
5.0
5.5
6.0
6.5
7.0
7.5
8.0
8.5
Field angle [ ]
38
Possible application of reflective filters (1)
This is a preliminary sketch of how reflective filters might be
incorporated. Problem is S-polarized reflection in blocking band. Next
slide shows solution.
MM
EFL=
D=
To radiator
f=
LIM
S
Scale
A.N
Heat pipe
39
Possible application of reflective filters (2)
The reflection of the S polarization component is large, but effect can be
minimized by crossing the direction of filters so S-pol becomes P-pol. The
second reflection would be to the side. This would simplify the heat pipe
also. Since detector is off to side, central telescope omitted (slide 41)
Second
reflector
Detector
and last
lens
MM
EFL=
D=
f=
LIM
S
Scale
A.N
Heat pipe
To radiator
40
4. Orbit, FOV, configuration,
Baffles
41
Orbit and FOV
 Sun synchronous polar orbit (inclined at °)
 For 720 km altitude,  =8°
± 23.4° + °
42
Stray Light
 Earth is illuminated by the sun up to 23.4° in winter/summer;
orbit plane inclined a further °
 In part of the orbit, this illuminated area, though outside the
FOV, will contribute stray light, most severely in the telescope
pointing nearest to this direction
 Stray light reaching detector will add to background level,
reduce sensitivity
 Baffles needed to exclude stray light, but cannot prevent it
completely
43
Stray light from Earth
Limiting angle to Earth (tangent depends on altitude- see table)
Earth
Altitude
Tangent
720
25.8
900
28.8
1100
31.5
1383
34.7
Field of
View
44
Baffles
 If stray light hits lens, impossible to suppress sufficiently (dashed red
line, baffle as dashed black line,)
 If baffle is long enough so stray light only hits baffle, it can be
suppressed by vanes, black coating (solid black line, solid red line)

Tangent
Length of baffle needed to
prevent direct light on lens
depends on angle 
45
Baffle Dimensions
  depends on altitude, field of view configuration
 Length of baffle to prevent direct stray light on lens given by
L  dlens / sin(  )
dlens = diameter of the lens
 The larger  the shorter the needed baffle
46
Telescope Configuration Options
9 telescopes,
36° x 36°
3.2% of sky
To use reflective
filters, centre unit
omitted, hence 8
8 telescopes,
2.8% of sky,
symmetrical
8 telescopes,
asymmetrical
8 telescopes,
asymmetrical,
2 rotated
47
Preliminary Mechanical Layout
Based on earlier 8 telescope asymetric option
48
Sun illumination on Earth
 Angle to illuminated part of earth
depends on position in orbit
 Longer arrow – lower angle
 Unit coloured orange may
be out of action due to
stray light when  too low
49
Stray Light effect on coverage
 If some telescopes are inoperative due to stray light during
part of the orbit for part of the year, the overall effect is small
 For example, one telescope out of action for 20 minutes on
each side of the orbit at peak sun inclination, for about 60% of
the year, this means 2*0.2*0.6/8 = 3% loss
 Since obscuration by the Milky Way will anyway reduce
observing time and may overlap with this loss, the effect will
be smaller.
50
Baffle length vs. Altitude
Minimum baffle length to avoid any loss of coverage due to stray light
Orbit Altitude [km]
Worst Angle  [°]
Baffle needed [mm]
720
5.9
1415
900
10.9
769
1100
16.4
512
1383
23.3
367
51
5. Status summary
52
System Performance
 Presently expected parameters:
# of telescopes
8
# of pixels per telescope
2048 x 2048
Pixel size
30 X 30 µm
F/#
2.4
Aperture diameter
120 mm
Integration time
300 secs
QE x transmittance
0.57 with reflective filter
PSF fraction on pixel
0.75
Bandwidth
50 nm
Background flux
0.08 ph/sec/cm²/sq. arcmin
53
Performance (continued)
 With transmissive filter, threshold signal photon rate is 0.011





ph/sec/cm²
At this level, detection distance is ~26Mparsec
At this distance, expected rate is ~2.4 per year
Proposed reflective filter has 95% transmittance although
higher background leak. Detection distance would be
40Mparsec and rate would increase by factor 2.66, i.e.
6.4/yr at 300 sec int. time
Wider bandwidth of reflective filter may increase rate even
more (needs further study)
Figure does not take into account obscuration by Milky Way
(possibly >20% of observation time)
54
Orbit Considerations
 Low earth orbits only ones practical, for cost reasons.
 Sensor must look in direction perpendicular to orbit plane to




allow 24 hour tracking after detection
Helio-synchronous orbit at LTDN 06:00 hrs is optimum
Due to earth axis inclination (±23.4°) and orbit inclination (8°
for 720km altitude, more for higher) will be some eclipse
periods unless altitude above 1400km
For same reason, some loss of energy due to angle to solar
panels. This may be minimized by bias
Baffle length needed is less at higher altitudes
55
Orbit Considerations
 Orbit stability is important, especially at lower altitudes. Drift




will cause increase in stray light, eclipse time
This will necessitate attitude control (momentum wheels,
magneto-torquers, etc.)
Line-of-sight stability needed: <50µrad in integration time to
limit smearing of image
Milky way will interfere with observation during part of the year
Proposed communication via SB-SAT constellation would give
gaps which increase with higher orbit. For this reason, lowest
orbit is preferable (720km)
56
SB-SAT Coverage
57
Signal processing
 To detect transients, image must be compared with a




reference image acquired before the transient occurred
Exact registration is essential. Shift of orbital plane (1°/day)
means interval cannot be long
Reference image must be average of at least 5 images to
reduce noise. If one image takes 5 minutes, this takes 25 min.
Bright objects will give high shot noise. Processing must be
conditional for such locations (e.g. wait for successive
detections)
Galaxies brighter than about mag.13 in one pixel (20’ x 20’)
could fall into this category but this seems most unlikely as
almost all galaxies are of lower intensity.
58
Uncertainties
 Background flux – not known if our value is correct
 Detector cooling – need to cool optics to reduce heat load on
detector to achieve required dark current.
59