Hyper Suprime-Cam
Design Review
March 2 2009

Contents
1 Introduction

5

2 Instrument requirements from Weak Lensing Sciences
2.1 Summary . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.2 Survey parameters . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.2.1 Number counts of galaxies usable for weak lensing analysis . . . . . .
2.3 Science goals of HSC WL survey . . . . . . . . . . . . . . . . . . . . . . . . .
2.3.1 Weak lensing power spectrum . . . . . . . . . . . . . . . . . . . . . .
2.3.2 Weak lensing cluster survey . . . . . . . . . . . . . . . . . . . . . . .
2.4 Photometric redshifts . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.4.1 Requirements for HSC WL Survey . . . . . . . . . . . . . . . . . . .
2.4.2 Simulations of photo-z Based on Our Method . . . . . . . . . . . . .
2.4.3 Estimating the photometric redshift error with clustering . . . . . . .

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3 Hyper Suprime-Cam System Deﬁnition
3.1 System Components . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2 Feasibility study of the HSC image quality . . . . . . . . . . . . . . . . . . .
3.2.1 Sensitivity Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2.2 Speciﬁcations of PFU mechanics . . . . . . . . . . . . . . . . . . . . .
3.2.3 Speciﬁcations for WFC and ADC . . . . . . . . . . . . . . . . . . . .
3.2.4 Speciﬁcations for ﬁlter . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2.5 Speciﬁcations for Camera mechanics and sensors . . . . . . . . . . . .
3.2.6 Summary . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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4 Wide Field Corrector (WFC)
4.1 WFC Requirements . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.2 WFC Design . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.2.1 Optics Layout . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.2.2 Designed Performance . . . . . . . . . . . . . . . . . . . . . . . . . .
4.3 Atmospheric Dispersion Corrector . . . . . . . . . . . . . . . . . . . . . . . .
4.4 Ghosting analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.5 Availability of glass blank . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.6 Design of Lens Barrel Assembly . . . . . . . . . . . . . . . . . . . . . . . . .
4.7 Impact of Lens Barrel Flexure on the Image Quality . . . . . . . . . . . . . .
4.8 Impact of the tilt and the displacement of CCD with respect to the lens . . .
4.9 Tolerance analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.10 Performance Veriﬁcation Test . . . . . . . . . . . . . . . . . . . . . . . . . .

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2

CONTENTS
4.11 Coating and the expected transmission . . . . . . . . . . . . . . . . . . . . .
4.12 Summary . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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5 HSC Camera Mechanics
5.1 CCD Dewar . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.1.1 Requirements . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.1.2 Mechanical Design . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.1.3 Thermal Design . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.1.4 Cooler . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.1.5 Vacuum Maintenance . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.1.6 Metrology of the Focal Plane Flatness . . . . . . . . . . . . . . . . .
5.1.7 Assembly Sequence . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.1.8 Prototyping and Development Schedule . . . . . . . . . . . . . . . . .
5.2 Shutter Assembly . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.2.1 Requirements . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.2.2 Mechanical Design . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.2.3 Prototyping . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.2.4 Development Schedule . . . . . . . . . . . . . . . . . . . . . . . . . .
5.3 Filter Exchange System . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.3.1 Requirements . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.3.2 Mechanical Design . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.3.3 FES Operation Sequence and Operation Time . . . . . . . . . . . . .
5.3.4 Control and Power System . . . . . . . . . . . . . . . . . . . . . . . .
5.3.5 Mounting of FES and Loading of Filters (Day-Time Operation) . . .
5.3.6 Development Schedule . . . . . . . . . . . . . . . . . . . . . . . . . .
5.4 Camera Mechanics System . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.4.1 Back Focus Conﬁguration . . . . . . . . . . . . . . . . . . . . . . . .
5.4.2 System Power Consumption . . . . . . . . . . . . . . . . . . . . . . .
5.4.3 Heat Exhaust System . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.4.4 Weight Estimate . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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6 Sensor
135
6.1 CCD Detector . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 135
6.1.1 Requirements for Hyper Suprime-Cam CCD . . . . . . . . . . . . . . 135
6.1.2 Hamamatsu fully depleted CCD . . . . . . . . . . . . . . . . . . . . . 136
6.1.3 Performance veriﬁcation . . . . . . . . . . . . . . . . . . . . . . . . . 143
6.1.4 Procurement and test plan . . . . . . . . . . . . . . . . . . . . . . . . 152
6.2 Front-end Electronics (FEE) and Back-end Electronics (BEE) . . . . . . . . 154
6.2.1 FEE and BEE . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 154
6.2.2 Front-end Electronics . . . . . . . . . . . . . . . . . . . . . . . . . . . 157
6.2.3 Back-end Electronics . . . . . . . . . . . . . . . . . . . . . . . . . . . 170
6.2.4 Schedule . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 181
185
7 Filter
7.1 Broad band ﬁlters for HSC . . . . . . . . . . . . . . . . . . . . . . . . . . . . 185
7.1.1 Requirements . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 185
7.1.2 Implementation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 188

CONTENTS
7.1.3 Prototyping at OCJ . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.1.4 Prototyping at BARR . . . . . . . . . . . . . . . . . . . . . . . . . .
7.1.5 Filter handling . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.1.6 Durability . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.1.7 Assembly of three substrates . . . . . . . . . . . . . . . . . . . . . . .
7.1.8 Filter glass holder assembly . . . . . . . . . . . . . . . . . . . . . . .
7.1.9 Ghost images . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.1.10 Prototyping and development status / schedule . . . . . . . . . . . .

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8 Telescope Interface
205
8.1 Telescope-HSC System Conﬁguration . . . . . . . . . . . . . . . . . . . . . . 205
8.2 Designs for Autoguider/Shack-Hartmann/Autofocus . . . . . . . . . . . . . . 207
8.2.1 Autoguider . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 207
8.2.2 Shack-Hartmann . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 209
8.2.3 Autofocus . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 210
8.3 Mechanical Modiﬁcation of Telescope . . . . . . . . . . . . . . . . . . . . . . 213
8.3.1 Inner hub . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 213
8.3.2 Clipping mechanism . . . . . . . . . . . . . . . . . . . . . . . . . . . 213
8.3.3 Cables . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 215
8.3.4 Telescope Balancer . . . . . . . . . . . . . . . . . . . . . . . . . . . . 215
8.4 Schedule for telescope modiﬁcation . . . . . . . . . . . . . . . . . . . . . . . 220
8.4.1 Tasks . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 223
8.4.2 Resources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 224
8.4.3 Estimated schedule . . . . . . . . . . . . . . . . . . . . . . . . . . . . 224
8.4.4 Impacts on operation of observatory . . . . . . . . . . . . . . . . . . 229
8.5 Safety Issues on Attachment of HSC to Telescope . . . . . . . . . . . . . . . 230
8.5.1 Concern on regular operations . . . . . . . . . . . . . . . . . . . . . . 230
8.5.2 Top unit exchange sequence . . . . . . . . . . . . . . . . . . . . . . . 230
8.5.3 Contact stress analysis . . . . . . . . . . . . . . . . . . . . . . . . . . 231
8.5.4 Preventive and post-fracture measures . . . . . . . . . . . . . . . . . 233
8.5.5 Proof experiment plan . . . . . . . . . . . . . . . . . . . . . . . . . . 234
9 Data Management and Analysis System
239
9.1 Suprime-Cam Data Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . 239
9.1.1 Current situation of Suprime-Cam data analysis . . . . . . . . . . . . 239
9.1.2 Problems in the current Suprime-Cam data analysis . . . . . . . . . . 240
9.2 Prototyping of HSC data analysis system . . . . . . . . . . . . . . . . . . . . 241
9.2.1 Approaches to overcome the problems in Suprime-Cam analysis . . . 241
9.2.2 Development steps . . . . . . . . . . . . . . . . . . . . . . . . . . . . 244
9.2.3 Goals of prototyping phase 1 and 2 . . . . . . . . . . . . . . . . . . . 245
9.2.4 Current status of the phase 1 prototype . . . . . . . . . . . . . . . . . 245
9.3 Key Issues of HSC Data Management and Analysis System . . . . . . . . . . 248
9.3.1 Diﬀerences between Suprime-Cam data and HSC data . . . . . . . . 248
9.3.2 Challenges to scale up Suprime-Cam prototyping system . . . . . . . 248
9.4 Goals of HSC Data Analysis System . . . . . . . . . . . . . . . . . . . . . . 249
9.4.1 Data products . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 249
9.4.2 Analysis time . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 250

4

CONTENTS
9.5

Summary . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 250

10 HSC Project Schedule

255

A Notes for observers

259

Chapter 1
Introduction
Dark energy has become one of the central puzzles in natural science since its discovery inlate 1990’s. Because of its tenuous distribution, only astronomical observation could probe
its existence. In order to understand the nature of dark energy, more precise observations
are crucial. However, the required survey speed, which is deﬁned by the product of the lightgathering power measured by the surface area A of the primary mirror and the camera ﬁeldof view Ω, is far beyond what can be provided by the existing astronomical facilities. Thisis a background of why next generation wide ﬁeld projects have been planned and carriedout, such as Dark Energy Survey (DES), Pan-STARRS and LSST.
Subaru Telescope has been playing a leading role in wide ﬁeld imaging astronomy for thelast decade. The prime focus provides φ 30 arcmin un-vignetted ﬁeld of view, which has
been unique among 8∼10 m class telescopes. Superb images better than 0.7 arcsec (FWHM)
are routinely obtained thanks to the precision optics and mechanism built by Japanese ﬁrms
(Figure 1.1). One example of Suprime-Cam images is shown in Figure 1.2(b) together with
an image from the Hubble Space Telescope. The narrowband image of Suprime-cam has anearly comparable in depth to the Hubble Deep Field, but cover more than a hundred times
the area.
It is interesting to note that the development project started in the early 1990’s, wellbefore the discovery of dark energy, and even before the discovery of Lyman break galaxies
by Steidel et al. which opened the door to the high redshift universe based on the opticalsurveys, and triggered the boom of optical wide ﬁeld surveys. In this sense, it was veryfarsighted of Profs. Kodaira (the ﬁrst Project Scientist of Subaru) and Okamura (PI ofSuprime-Cam) to make a decision to implement a prime focus on Subaru Telescope.
Hyper Suprime-Cam (HSC) is a proposed upgrade of Suprime-Cam. In order to makethe survey speed even faster, we plan to expand the ﬁeld size by nearly 10 times while
maintaining the equivalent image quality as that of Suprime-Cam. The design goal of HSC
is summarized in Table 1.1. Compared with other camera and survey projects, such as
DES and Pan-STARRS, one of our advantage is that Subaru is the only telescope that has
already demonstrated high image qualities whose meadian seeing is better than 0.7 arcsec(FWHM). This is especially crucial for weak lensing survey. Takeing a full advantages ofthis opportunity to do even more exciting things, we are keen to complete the instrument
and start the unique and legacy survey in a timely manner.
56

CHAPTER 1. INTRODUCTION
Field of View

1.5 degree diameter

Instrument PSF size
Pixel scale
System Throughput

Minimum Exposure time
Min. interval of Exposures≤ 0”.4 for g’, r’, i’, z’≤ 0”.2 /pix
≥ 50 % for g’≥ 65 % for r’≥ 65 % for i’≥ 40 % for z’≥ 20 % for Y
2 sec (1 sec goal)
20 sec (15 sec goal)

Min. number of ﬁlter holders
Filter Exchange Time4
< 10 mininutes

Vignetting allowed up to 25 % at the edge
Dead area (CCD gap) fraction ≤ 5 %FWHM∗ , Telescope Elevation > 30 deg∗∗
PM×WFC×Filter×CCD
at the center of the ﬁeld

Time accuracy ≤ ± 1 %Including CCD readout and wipe
pointing change

Table 1.1: Primary speciﬁcation of Hyper Suprime-Cam. ∗ D80 = 1.56× FWHM for gaussian
distribution. ∗∗ Observing record tells that 95 % of Suprime-Cam observation is made between
30 < EL < 85
Purpose of this book and the contributors
This book summarizes the scientiﬁc and techinical studies regarding HSC, and is prepared
for a Design Review organized by Prof. Miyama, Director General of National Astronomical
Observatory of Japan. The following is a list of the HSC project members and Subaru
telescope members who contribute to the review booklet.
H. Aihara
Y. DoiH. Furusawa
T. Hamana
S. Ichikawa
Y. Ishizuka
R. ItoN. Katayama
Y. KamataH. Karoji
S. Kawanomoto
Y. Komiyama
H. Miyatake
S. Miyazaki
T. Morokuma
H. Nakaya
K. Nariai

Univ. of Tokyo
NAOJ
NAOJ
NAOJ
NAOJ
NAOJ
KEK
KEK
NAOJ
NAOJ
NAOJ
NAOJ
Univ. of Tokyo
NAOJ
NAOJ
NAOJ
NAOJ

H. Nishioka
A. Nishizawa
Y. Obuchi
N. Okabe
M. Takada
T. Takata
M. Tanaka
Y. Tanaka
D. Tomono
T. Uchida
K. Umetsu
F. Uraguchi
Y. Utsumi
S. Wang
N. Yasuda
K. Yamamoto

ASIAA
NAOJ
NAOJ
Tohoku Univ.
Univ. of Tokyo
NAOJ
KEK
NAOJ
NAOJ (Subaru Telescope)
Univ. of Tokyo, KEK
ASIAA
NAOJ
NAOJ
ASIAA
Univ. of Tokyo
Hiroshima Univ

7

Figure 1.1: (a) Focal plane of Suprime-Cam which consists of ten MIT/LL CCID-20 CCDs
mounted in the vacuum dewar. The physical size of the mosaic is 15 cm × 12 cm andthe longer side corresponds to 30 arcmin. (b) Subaru prime focus corrector lens designed
and fabricated by Canon. The diameter of the ﬁrst lens is roughly 50 cm. (c) Prime focus
structure, designed and built by Mitsubishi Electric, which includes an instrument rotator,
a hexapod with precise actuators for attitude control and a cable wrapper.

Figure 1.2: (a) The Hubble Deep Field imaged with the WFPC2 on HST through theF814W continuum ﬁlter. (b) The same region imaged from the ground through a narrow
band NB816 ﬁlter using Suprime-Cam. (Courtesy of E. Hu from their review in Nature).

Chapter 2
Instrument requirements from Weak
Lensing Sciences
2.1

Summary

In this chapter, we describe a plan and feasibility of a HSC wide ﬁeld cosmology survey andits science cases focusing mainly on weak lensing related sciences. This proposal is made by
HSC weak lensing working group based on their feasibility study of HSC wide ﬁeld cosmology
survey.
We propose a HSC wide ﬁeld survey for the cosmology which is designed to provide afundamental data set for various areas of cosmology-related researches in the next decadeby making the best use of HSC/Subaru abilities. It is designed to cover a cosmologicalvolume and to reach the limitting magnitude of at least one magnitude fainter than otheron-going/planed wide ﬁeld surveys in the HSC era conducted with 4m class telescopes (e.g.,
the CFHT Legacy Survey and the Dark Energy Survey). The basic survey parameters are
as follows.
• Survey area is 2000 square degrees to be observed with four passbands (g, r, i, z).
>• The survey depths are ∼26 AB magnitude in g-, r-, and i-band and ∼ 25 AB magnitude in z-band. The required exposure time for each pointing are 15 minutes for g-,
r-, and i-band and 20 minutes for z-band (thus 65 minutes in total).

The main focus of the HSC wide ﬁeld survey is weak lensing related sciences, so we exchangeably refer to it as HSC weak lensing (WL) survey. The HSC WL survey aims at twogoals: The ﬁrst is to explore a cause of the accelerating expansion of the universe. HSC WL
survey will allows an accurate tomographic measurement of the weak lensing shear power
spectrum from which we will be able to place a tight constraint on the dark energy equationof state parameter (w) and alternatively to test the derivation of the gravity from Einstein’s general relativity (§2.2). The second is to generate, for the ﬁrst time, a large catalog(N ∼ O(105 )) of mass selected clusters of galaxies from weak lensing mass map. The cluster
catalog will provide a valuable way to exploring the cosmic structure formation. Moreover,the catalog, combined with cluster samples generated by surveys in other wavelength (e.g.,the optical, X-ray and radio), will allows to probe the formation and evolution of both thebaryon and dark matter components in the clusters and galaxies (See section2.3.2).
9

10CHAPTER 2. INSTRUMENT REQUIREMENTS FROM WEAK LENSING SCIENCES
To archive the science goals, an accurate photometric redshift (photo-z) is essential. We
present results of our photo-z simulations to examine the performance of the photo-z obtained by griz passbands data. We ﬁnd that with only the four bands, it is possible toarchive the photo-z accuracy required from the HSC WL study if photo-z outliers can be
successfully excluded. In order to establish eﬃcient and robust ways of removing the catastrophic errors, we will improve our photo-z method in three ways: (1) implementing theempirical training-set approaches, (2) we will introduce reasonable priors, and (3) we will
introduce a “reliability parameter” of resulting photo-z (see §2.4 and references therein for
details). We will also utilize the galaxy clustering to self-calibrate the photometric redshifterror (§2.4.2).
The feasibility of sciences goals mentioned above also relies on instrument speciﬁcation of
HSC. Here we summarize requirements from weak lensing sciences.
1. The most critical requirement is for the size of the ﬁeld-of-view. In order to obtaina wide ﬁeld imaging data demanded by science goals (see the following sections fordetails) within a reasonable observation time, at least 1.8 square degrees or 1.5 degrees
φ is required (see §2.2). Also for an eﬃcient survey, a short readout times (a few tens
seconds) is required.
2. Since a good image quality is crucial for the weak lensing analysis, natural seeing
limited imaging should be realized in most of the observing time.
3. In order to handle a huge data quickly and stably, a full automatic data reduction
pipeline is required.

2.2

Survey parameters

The HSC weak lensing (WL) survey described below is designed to provide a fundamental
data set for various areas of cosmology-related researches in the next decade by making thebest use of HSC/Subaru abilities, and its primary target is weak lensing sciences describedin the following sections. Principal guidelines for the HSC WL survey design are as follows:• It is especially aimed at reaching about one order fainter ﬂux than other on-going/planed
wide ﬁeld surveys in the HSC era conducted with 4m class telescopes (e.g., the CFHT
Legacy Survey (CFHTLS)/CFHT1 and the Dark Energy Survey(DES)/CTIO2 ).
• It should cover a cosmological volume, to be more speciﬁc, it should be large enough tosample a statistically signiﬁcant number of largest cosmic structures such like superclusters and voids.
The former leads to the survey depth of ∼26th AB magnitude in g 0 -, r0 -, and i0 -band and∼25th AB magnitude in z 0 -band. This makes a big diﬀerence to the galaxy number density
usable for a weak lensing analysis; a typical number density of CFHTLS/DES is ng ∼ 13
arcmin−2 (e.g., Fu et al. 2008) whereas one of HSC WL survey is about three times larger,
ng ∼ 35 arcmin−2 (see §2.2.1). This especially aﬀects two points: The ﬁrst is that the
HSC WL survey can probe the dark energy equation of state at higher redshift than the12

http://www.cfht.hawaii.edu/Science/CFHTLS
https://www.darkenergysurvey.org

2.2. SURVEY PARAMETERS

11

Figure 2.1: Mass functions of dark matter halos predicted from the extension of the PressSchecter model by Sheth & Tormen (1999) in a ﬂat ΛCDM cosmology (ΩΛ = 0.7) for diﬀerent
redshifts (z = 0, 0.5, 1.0 and 2.0).
Table 2.1: Proposed survey parameters. SN > 5 detection of point source with 200 diameter
aperture under 000 .8 seeing.
Filtergriz
Depth [AB mag]
26.4 26.1 25.8 24.8
Exposure time [min] 15
15
15
20
shallower surveys (see §2.3.1 and Figure 2.7). The second is mass and redshift ranges ofgalaxy clusters to be detectable with weak lensing technique. In fact, the HSC WL survey
will be able to locate about four times less massive clusters than the shallower surveys (see§2.3.2 and Figure 2.11). Thus the HSC WL survey alone can provide a sizable (N ∼ 105 )catalog of mass selected typical mass clusters, which will provide valuable synergy with a SZ
or X-ray cluster catalog.
On the other hand, concerning the latter, in order to estimate a relevant cosmologicalvolume we consider expected number counts of the most massive clusters of galaxies. Below,
we evaluate a required survey area to sample more than one hundred most massive clusters.
To do this, we show in Figure 2.1 the CDM prediction of the mass function of the dark
matter halos. It is found from the ﬁgure that the number density of dark matter halowith mass M ∼ 1015 h−1 M¯ at redshift range of our interest (0.1 < z < 0.7) is of orderO(102 )/(h−3 Gpc3 ). Therefore the relevant survey volume should be more than 1 h−3 Gpc3 .
Since the comoving light-cone volume to z = 1(2) is ∼ 1.2(4.6)[h−3 Gpc3 ]/1000 deg2 , morethan 1000 deg2 ﬁeld is required.
Here we evaluate a required survey time. To do so, we suppose a ﬁducial survey area of2000 deg 2 and the HSC’s ﬁeld-of-view of 1.5 deg φ. The required exposure times for eachﬁlter are summarized in Table 2.1. A required total exposure time is 1200 hours or 150
nights.

12CHAPTER 2. INSTRUMENT REQUIREMENTS FROM WEAK LENSING SCIENCES

Figure 2.2: Candidate ﬁelds of HSC weak lensing survey are shown. The orange regions
show the survey ﬁelds of the Large Area Survey (LAS) of the UKIRT Infrared Deep Sky
Survey (UKIDSS), the purple region shows the VIKING-N survey ﬁeld, the yellow regions
show the CFHTLS wide ﬁelds including the XMM-LSS ﬁeld and the UKIDSS-DXS (Deep
Extragalactic Survey) ﬁelds, and red region shows the SDSS strip-82 ﬁeld (also the UKIDSSLAS survey ﬁeld). The hatched blue regions (30 < |b| < 70deg) show preferable regions forweak lensing survey, where the number density of stars (that are used to measure shapes ofPSF) is appropriate.

2.2. SURVEY PARAMETERS

13

Finally, let us devote a little more space to discuss a survey region. Although the HSC
WL survey is basically designed as the stand-alone program, it is clear that the more multiwavelength surveys it coincides with, the more valuable synergy we may have.
Near-IR and u-band data will improve the accuracy of photometric redshift of source
galaxies and galaxy clusters, and thus will provide more accurate constraints on cosmologicalparameters. There are two strong candidate ﬁelds: One is the VIKING-N ﬁeld3 (780 deg2
in the equator from R.A. 10h to 16h and between Dec. -5 and 3 deg; J = 22.1, H = 21.5,
Ks = 21.2 (in AB magnitude) shown in Figure 2.2 as the purple region). The other is theLarge Area Survey (LAS) of the UKIRT Infrared Deep Sky Survey (UKIDSS)4 (1900 deg2
in the equator, 1900 deg2 in northern sky, and 200 deg2 in the equator; Y = 20.5, J = 19.7,
H = 18.8, K = 18.4 in Vega magnitude, shown in Figure 2.2 as the orange regions). The
CFHT Legacy wide survey takes the deep u-band image, thus it’s survey ﬁelds (4 ﬁelds
shown by yellow squares in Figure 2.2, the total area is 170 degree2 ) are the candidate of apart of the HSC WL survey area.
Also a massive spectroscopic catalog will be helpful to calibrate and thus to improve thephotometric redshift technique. The SDSS ﬁelds, especially ﬁelds where the BOSS5 projects
to be carried out, are strong candidates of the survey area.
A possible synergy with deep X-ray surveys (e.g. eROSITA, a document is available fromhttp://www.mpe.mpg.de/erosita/MDD-6.pdf and Sunyaev-Zel’dovich eﬀect(SZE) surveys(e.g., ACT) are of great interest, because those provide with informations on physical stateof the baryonic matter in clusters of galaxies. SDSS stripe-82 ﬁeld (the red region in Figure2.2) is one of candidate survey ﬁelds of the ACT SZE survey6 . Also, areas where deep
X-ray image taken suck like XMM-Newton Large-scale Structure Survey (XMM-LSS) arestrong candidate as it will provide valuable data for probing physical properties of clustersof galaxies.

2.2.1

Number counts of galaxies usable for weak lensing analysis

Since the statistical noise of the weak lensing mass map and weak lensing power spectrum√
mainly comes from the intrinsic galaxy shape, which scales as σe ∝ σint / ng , where σint
is the root-mean-square of the intrinsic galaxy ellipticities and ng is the number density ofgalaxies used for a weak lensing analysis. Therefore, an accurate knowledge of the relation
between the number counts of galaxies usable for weak lensing analysis and the exposure
time is of essential importance to optimize a survey design. For example, in the case of
the weak lensing cluster survey, the higher the galaxy density becomes, the higher the S/Nof weak lensing signal from a cluster becomes, and thus the more clusters one can locate.
Also, a deeper survey allows us to explore structures at higher redshifts. However, given alimited observational time, a deeper imaging by a longer exposure time naturally results ina narrower survey area. Thus in order to make theoretical predictions of products from aweak lensing survey and to make the best survey design, a knowledge of a correct relation
between the galaxy number density and the exposure time is very important. However, it
is diﬃcult to make a reasonably accurate model of this, because in a weak lensing analysis,
galaxies smaller than the point spread function (PSF) are discarded as their shapes are3

http://www.eso.org/sci/observing/policies/PublicSurveys/sciencePublicSurveys.html
http://www.ukidss.org/
5
http://www.sdss3.org/cosmology.php
6
http://www.physics.princeton.edu/act/index.html
4

14CHAPTER 2. INSTRUMENT REQUIREMENTS FROM WEAK LENSING SCIENCES

Figure 2.3: The number counts of galaxies usable for a weak lensing analysis as a functionof the exposure time. The Suprima-Cam i0 -band data is used. Diﬀerent symbols are fordiﬀerent seeing conditions (0.5, 0.7 and 0.9 arcsec from upper to lower).
smeared by the PSF and are very poorly measured. Also, for the same reason, we can not
use the standard galaxy number counts. Instead, we used actual Suprime-Cam archive data
to evaluate the relation. We used deep i0 -band image taken by the Subaru-XMM-Deepsurvey project7 . By combining many short time exposure time images, we generated 5, 15,30, 45 and 60 minutes integrated images for three seeing conditions of ∼ 0.5, 0.7 and 0.9
arcsec. Then we did object detections with IMCAT as in the same manner as was donein a usual weak lensing analysis (e.g., Miyazaki et al. 2007). The galaxy counts obtained
are shown in Figure 2.3. We found that the galaxy counts are ﬁtted reasonably well by thefollowing simple power-law model:−1ng ∝ t0.3(2.1)
exp s ,
where s is the seeing size measured by the FWHM of point sources (stars). It is thus foundthat the number counts is very sensitive to the seeing condition. This makes the Subaru oneof the best ground-based telescopes for weak lensing survey because its good seeing condition(the median seeing is about 0.6 arcsec, Miyazaki et al. 2002). The empirical relation, eq.
(2.1) is used to make theoretical predictions in the following sections.

2.3

Science goals of HSC WL survey

2.3.1

Weak lensing power spectrum

Weak gravitational lensing due to cosmological large-scale structures, the so-called cosmic
shear, causes images of distant galaxies to be distorted. The shearing eﬀect is mostly too
weak to measure from a single galaxy, but is indeed measurable by summing up the correlated
distortions over a suﬃciently large number of galaxies. The method most conventionally used7

http://www.naoj.org/Science/SubaruProject/SDS/

2.3. SCIENCE GOALS OF HSC WL SURVEY

15

in the literature is the two-point correlation function of the cosmic shear ﬁelds (e.g., Hamana,
Miyazaki et al. 2003). The precise measurement of cosmic shear correlations requires a huge
number of well-resolved galaxies usable for the lensing analysis as well as high image quality
in order to discriminate the lensing signal from contamination due to atmospheric seeing
eﬀect and telescope distortion. For this reason, the HSC survey proposed in § 2.2 oﬀers an
ideal, unique opportunity to measure the cosmic shear signals with unprecedented precision
thanks to its widest ﬁeld-of-view, among other 8m class telescopes, and the excellent image
quality of the Subaru Telescope.
The cosmic shear correlation function measures the total matter power spectrum weighted
with a function that depends on a combination of the angular diameter distances for thesource-lens-observer system. The redshift evolution of the matter power spectrum as well as
the angular diameter distance combination are both sensitive to properties of dark energy.
Utilizing this sensitivity, the precise cosmic shear measurement can be used to determine
whether the cosmic accelerated expansion observed today is caused by the cosmologicalconstant or a more generalized energy component, dark energy. An even more ambitious
use of the cosmic shear is to explore an alternative possibility that the cosmic accelerationis, rather than dark energy, as a consequence of a modiﬁcation of Einstein’s gravity on
cosmological distance scales, because a modiﬁcation of gravity generally predicts diﬀerent
growth history of mass clustering even for a ﬁxed cosmic expansion rate (Yamamoto et al.
2006, 2007). In particular, the gravitational lensing based methods, combined with thegeometrical methods such as SNe and baryon acoustic oscillation experiments, open up sucha window to test the Einstein gravity on cosmological scales, as recommended in the Reportof Dark Energy Task Force8 .
The multi-color data set of the HSC WL survey can be used to estimate the distances ofgalaxies, i.e. photometric redshifts (hereafter photo-z), by utilizing the fact that galaxies
display characteristic photometric features depending on the redshift and type of galaxies(also see § 2.4 for the further discussion). This additional information is extremely valuable
in that it allows us to recover redshift evolution of the matter distribution and the lensing
geometrical function from the measured lensing signals. To be more explicit, measuring thecosmic shear correlations as a function of photo-z’s and separation angles of galaxies greatly
improves the ability of the lensing survey for constraining the geometry of the universe as
well as the growth history of mass clustering, which can be in turn used to constrain thedark energy properties or test the Einstein’s gravity (Takada & Jain 2004; Takada & White
2005). This method is now dubbed lensing tomography and has emerged as one of the mostpromising to constrain cosmology if the inherent systematics are well under control.
In addition to the statistical lensing measurement described above, measuring the cosmic
shear enables us to directly map out the two-dimensional distribution of total matter (darkmatter plus baryon in ratio approximately 10:1) projected along the line-of-sight, irrespectiveof uncertainties associated with galaxy bias. A high peak in the reconstructed mass map will
provide a strong candidate for a massive cluster of galaxies (Miyazaki, Hamana et al. 2002;
Miyazaki, Hamana et al. 2007; also see § 2.3.2 for the details). In addition, the multi-color
data set itself will help discriminate real clusters from false peaks due to the intrinsic galaxy
shapes or the projection eﬀect, because member galaxies in clusters provide secure photo-z
estimates such as prominent red sequence of member galaxies. A cluster catalog can thus beconstructed from the HSC WL survey. Further, the survey region may be overlapped with
8

The documents are available from http://www.nsf.gov/mps/ast/detf.jsp.

16CHAPTER 2. INSTRUMENT REQUIREMENTS FROM WEAK LENSING SCIENCES

Figure 2.4: The cosmic shear power spectra for galaxy distribution divided in three redshift slices,0 ≤ z1 ≤ 0.6, 0.6 ≤ z2 ≤ 1, and z3 ≥ 1. The bold solid curves show the auto-spectra of 3 redshiftbins for a ΛCDM model, while the thin curves are the results for a model with the dark energy
equation of state w = −0.9. The boxes around the bold curves show the expected measurement
error due to the sample variance and the intrinsic galaxy shapes, for our proposing HSC WL surveythat is parametrized by survey area Ωs = 2, 000 deg2 , the average number density of galaxies usable
for lensing analysis, n̄g = 30 arcmin−2 , and the rms intrinsic elipticities per component, σ² = 0.22.

other cluster surveys in diﬀerent wavelengths, the X-rays and/or the Sunyaev-Zel’dovich
eﬀect, and the combination will strengthen the reliability of the cluster catalog. It is wellknown that number count of clusters is a powerful probe of cosmology, through the sensitivityof the halo mass function and the comoving volume element to cosmological parameters (e.g.
Kitayama & Suto 1997). Takada & Bridle (2007) have recently demonstrated that combining
the number counts of clusters and the cosmic shear correlation can be a powerful way to
improve cosmological parameter determination, even if the two observables are drawn fromthe same survey region. We will below describe an expected performance of the HSC WL
survey for constraining dark energy parameters from the combined measurement.
2.3.1.1

Feasibility and expected results

The Fourier transformed counterpart of the cosmic shear correlation function is the angularpower spectrum. Since the shear ﬁeld has the spin-2 ﬁeld properties, the measured shear
power spectra are divided into two independent spectra of curl-free (E) and curl (B) modes.
The consequence of weak lensing caused by a scalar gravitational potential is that only the
E-mode spectrum is allowed, therefore the magnitude and scale dependence of B-mode in2.3. SCIENCE GOALS OF HSC WL SURVEY

17

the data can serve as a useful monitor of remaining systematics.
The bold solid curves in Figure 2.4 show the expected lensing power spectra for 3 redshift
bin tomography for a concordance ΛCDM model that is consistent with the recent WMAP
result (Komatsu et al. 2008). We here assumed survey parameters that approximately
correspond to the HSC WL survey: the survey area Ωs = 2, 000 deg2 , the redshift distributionof galaxies peaks at z = 0.6 and has mean redshift hzi = 0.9, and the average number density
ng = 30 arcmin−2 . Note that the tomographic redshift bins are simply deﬁned so that eachbin contains roughly equal number density of galaxies; 0 ≤ z1 ≤ 0.6, 0.6 ≤ z2 ≤ 1 andz3 ≥ 19 . The relative shifts in the amplitudes and angular-scale dependences of these
spectra are sensitive to the dark energy properties. The boxes around each curve show the
expected measurement error due to sample variance and intrinsic galaxy shapes. It canbe seen that the HSC WL survey allows a signiﬁcant detection of the lensing spectra ineach redshift bin and over a wide range of multipoles, especially around l ∼ 103 or ∼ 10
arcminute scales. Increasing the dark energy equation state from w = −1 to w = −0.9
causes the lensing power spectrum amplitudes to be lowered, because the change slows down
the growth of mass clustering at relevant redshifts as well as lowers the lensing eﬃciency due
to the decreases in the angular diameter distances. This can be explicitly demonstrated by
the thin solid curves that show the spectra for a model with w = −0.9 but other parameters
being ﬁxed. The diﬀerence between the bold and thin solid curves, compared to the errors,
implies that the two models with |∆w| = 0.1 may be easily discriminated. However, torealize a genuine ability of the HSC WL survey for constraining the dark energy parameters,
we need to properly take into account degeneracies between all cosmological parameters thataﬀect the lensing power spectra, because a combination of the parameters may mimic thedark energy dependence. A realistic parameter forecast will be presented below.
On the other hand, Figure 2.5 demonstrates the expected redshift distribution of clusters
that are selected from lensing mass peaks reconstructed from the same HSC WL survey. The
dashed, solid and dotted curves show the results for the detection thresholds (S/N )cluster =>6, 8 and 10, respectively. In practice, such a high detection threshold of (S/N )cluster ∼5needs to be employed in order to avoid contamination of false peaks due to the intrinsic
galaxy shapes and the projection eﬀect of structures unassociated with a cluster (Hamana
et al. 2004). Compared to the mass-selected cluster sample, the lensing-selected cluster
sample has a narrower redshift coverage reﬂecting the redshift dependence of the lensing
eﬃciency for typical source redshift z ∼ 1. A change of w from w = −1 to w = −0.9
leads to a decrease in the number density of clusters as shown by the thin solid curve. Thiseﬀect arises partially from the decrease in the comoving volume element and the halo massfunction, and an additional contribution is from the lowered cluster lensing signal.
We now quantify how the HSC WL survey will improve our understanding of dark energy
properties when using a combined measurement of the lensing power spectrum tomography
and the cluster counts. Figure 2.6 shows the expected error ellipses in dark energy parameter
space. The parameter forecasts shown assesses how well the given observables can distinguish the true (“ﬁducial”) cosmological model from other models, therefore the forecasts
depend on choices of the ﬁducial model and free parameters and also vary with priors used.
We assume a current concordance ΛCDM model for the ﬁducial cosmology and employ 9
cosmological parameters in total including time evolution of the dark energy equation of
9

For lensing tomography with nz redshift bins, there are nz (nz + 1)/2 auto- and cross-spectra constructedusing galaxy pairs sitting into same or diﬀerent redshift bins.

18CHAPTER 2. INSTRUMENT REQUIREMENTS FROM WEAK LENSING SCIENCES

Figure 2.5: The average number density of lensing-selected cluster sample, per unit square arcminute and per unit redshift interval, for the ΛCDM model. Here clusters having a lensing signal
greater than a given threshold are included in the sample. The dashed, solid and dotted curves are
the results for the detection thresholds (S/N )cluster = 6, 8 and 10, respectively. For comparison, thetwo dot-dashed curves show the number density for the mass-selected cluster sample with masses
M/M¯ ≥ 5, 10 × 1014 . Increasing the dark energy equation of state from w = −1 to w = −0.9
lowers the number density as shown by the thin solid curve, compared to the bold-solid curve.

state parametrized asw(a) = w0 + wa (1 − a),(2.2)(see Takada & Bridle 2007 for the details). The errors shown in Figure 2.6 include marginalization over other parameters. Note that we assume the CMB priors expected from thePlanck satellite mission. However, note that the errors on w0 and wa are enlarged only by∼ 10% even if very conservative priors, σ(ln(Ωb h2 )) = 0.05 and σ(ln(Ωm h2 )) = 0.1, (no prioronto other parameters) are employed instead of the Planck priors.
Comparing the dashed and intermediate solid contours, one can clearly ﬁnd that the use oflensing tomography helps signiﬁcantly break the parameter degeneracies, yielding an isolated
error ellipse rather than a perfect degenerated ellipse in the dark energy parameter space.
Adding the number counts of lensing-selected clusters with detection threshold (S/N )cluster =6 further shrinks the error ellipses due to the complementarity between the lensing power
spectrum tomography and the cluster counts in cosmological parameter determination. As
a result, the projected 68% uncertainties σ(w0 ) ≈ 0.14 and σ(wa ) ≈ 0.39 may be attained.
For the constant dark energy equation of state, σ(wconst ) ≈ 0.048.
We should bear in mind that, to exclude the cosmological constant observationally, it is2.3. SCIENCE GOALS OF HSC WL SURVEY

19

Figure 2.6: An expected performance of the HSC WL survey for constraining the nature ofdark energy parametrized by three parameters, the present-day energy density parameterΩde and the dark energy equation of state parameters w0 and wa for the parametrizationw(a) = w0 + wa (1 − a). The dashed contour shows the 68% C.L. error ellipses for the lensing
measurement without tomography, while the intermediate solid contour shows the result for
the three-redshift bin tomography as shown in Figure 2.4. Note that we include all the autoand cross-spectra constructed from the three redshift bins and consider multipole range 50 ≤l ≤ 3000. Further, the innermost ﬁlled contour is the result when the lensing tomography
is combined with the number counts of lensing-selected clusters shown in Figure 2.5, where
the cluster counts are subdivided into 10 redshift bins over redshift 0.05 ≤ z ≤ 1. The
constraints are marginalized over uncertainties on other parameters (9 parameters are used
in total), assuming the CMB priors.

suﬃcient to ﬁnd w 6= 1 at a particular single redshift. Therefore, it would be useful toknow the pivot redshift at which the dark energy equation of state is best constrained with
the HSC survey. To study this, we use an alternative parameterizaion of the dark energy
equation of state given by w(a) = wp + wa (ap − a) instead of Equation (2.2), where wp , wa
and ap are treated as free parameters. Figure 2.7 shows the marginalized errors on wp asa function of zp = 1/ap − 1. One can ﬁnd that the best-constrained error σ(wp ) ∼ 0.05 isavailable at the pivot redshift zp ∼ 0.55. It would be also worth noting that the errors on
wp and wa are uncorrelated at this pivot redshift. For comparison, we also show the result
for a shallower survey, here speciﬁcally a survey like the Dark Energy Survey (DES), whichprobes the Universe with medium redshift hzs i = 0.7 and has the similar-level precisionof constraining the dark energy equation of state; in other words, HSC and DES are bothcategorized as the Stage III surveys in DETF. Such shallow surveys are sensitive to darkenergy at a slightly lower redshift z ∼ 0.5. More importantly, a deeper survey generally
has a wider window of redshifts to probe dark energy; the HSC survey has a sensitivity of20CHAPTER 2. INSTRUMENT REQUIREMENTS FROM WEAK LENSING SCIENCES

Figure 2.7: The marginalized error on wp for the HSC survey (solid curve) as a function of
zp (= 1/ap − 1), for parameterization of the dark energy equation of state given by w(z) =wp + wa (ap − a), instead of Eq. (2.2). The best-constraint error, σ(wp ) ' 0.05 is available atthe pivot redshift zp ' 0.55 for the HSC survey. For comparison, the results for a shallower
survey like the DES survey is also shown. The HSC survey has a wider window of redshifts
for constraining the dark energy equation of state than the DES has.

<<<σ(w) ∼
0.07 over a range of 0.4 ∼z∼0.8, while the DES-like survey has a much narrower
window.

2.3.1.2

Requirements and prospects for systematics

For the parameter forecasts shown in Figure 2.6, we ignored the impacts of systematic errors
inherent in weak lensing measurement and cluster counts on parameter determination forsimplicity. Here we comment on requirements on a level with which the systematic errors
need to be understood in order to achieve the full potential of the HSC WL survey. Notethat the issues discussed below are still in working progress, and have yet to be deﬁnitely
clariﬁed.
Photometric redshift errors
For weak lensing, most problematic systematics are caused by biases in photo-z errors andthe shape measurement of galaxies. For the requirement on the photo-z errors, following themethod developed in Huterer et al. (2006), we ﬁnd that a mean redshift of each tomographic
redshift bin needs to be known with ∼ 1% accuracy in order not to have 100% bias in the darkenergy parameter constraints. To achieve the % level accuracy with 4 or 5 broad-band colorsof the HSC WL survey seems still challenging (see § 2.4 for more detailed discussion). The2.3. SCIENCE GOALS OF HSC WL SURVEY

21

statistical random errors on photo-z are not problematic, because the errors are expected tobe typically σ(z) ∼ 0.1 for individual galaxies and ∼ 108 galaxies are available in each redshift
bin to suﬃciently reduce the error contribution to the mean√redshift after the statistical
average; the residual error is roughly estimated as δz ∼ 0.1/ 108 ∼ 10−5 . Hence, it is of
critical importance to remove the contamination caused by catastrophic photo-z outliers.
The multi-color imaging data set alone may be insuﬃcient to control this systematics, anda training spectroscopic (internal or external) galaxy sample may be needed for the photoz error calibration. Important to note for this issue is, for the planned HSC WL survey,
the lensing measurement is not shot noise limited as long as only the spectra of multipole<
range l ∼3, 000 are included – a signiﬁcant fraction of galaxies, say up to one third, canbe discarded if their redshifts are ill-deﬁned without severely degrading parameter accuracy(e.g., see Jain, Connolly & Takada 2007). Thus the interesting issue to be addressed iswhether it is possible to deﬁne a sub-sample of galaxies that have secure photo-z estimates
well suited for lensing tomography, using multi-color data set alone and/or combining withtraining spectroscopic data set.
Shape measurement errors
For the shape measurement uncertainty, a bias in the shear amplitude needs to be knownwith a few 10−3 accuracy. For this systematics, there is an important development that has
recently been made: the international collaboration project, the Shear TEsting Programme(STEP), aimed to improve the accuracy and reliability of existing weak lensing measurement
techniques (Massey et al. 2007). Encouragingly, some of most accurate methods are not so
far from the required accuracy for the HSC WL survey, so this systematics may not be alimiting factor in ∼5 years when the HSC WL survey will actually start.
Uncertainty in modeling WL observables
Another concern is inherent in uncertainties in predicting the nonlinear clustering of structure formation to which the weak lensing observables are sensitive. The sensitivity of weaklensing surveys peaks at scales of about 10 arcminutes on the sky (see Figure 2.4), corresponding to intervening structures of ∼ 1h−1 Mpc at redshift ∼ 0.5, which is approximately
where the lensing eﬃciency of distant galaxies peaks. Therefore, since the scales probed
are in the nonlinear clustering regime, accurate calibration of the nonlinear matter power
spectrum Pm (k) based on N -body simulations is necessary to ensure that the cosmologicalparameter determination from HSC survey is not compromised10 .
Following the method developed in Huterer & Takada (2006), we have found that calibration of Pm (k) at scales around k ∼ 1Mpc−1 needs to be better than ∼ 3% for the HSC WL
survey. Current accuracy in the matter power spectrum up to the mildly nonlinear scales
is about 5-10% (Vale & White 2004; Heitman et al. 2004), which is not suﬃciently small,but the required accuracy is nearly within reach with current simulations. We will therefore
need to resort to a full suit of N -body simulations, interpolating on a grid of cosmologicalmodels. First attempts in such an approach has been recently undertaken, mainly led by thegroup at Nagoya University (Takahashi et al. 2007), and the required calibration of P (k)is very expected to reach until the HSC survey is coming on line, although the number ofmodels that is necessary to be run might be large (and therefore, perhaps, building a new
10

Unfortunately, one cannot aﬀord to restrict the measurements to the linear regime for which there are>
accurate analytic predictions and which correspond to scales ∼1 degrees, since the information from large
scales has only poorly constraining power due to cosmic variance.

22CHAPTER 2. INSTRUMENT REQUIREMENTS FROM WEAK LENSING SCIENCES
set of more ﬂexible and accurate ﬁtting formulae may be needed).
Mass-observable relation for clusters
For the cluster counts, a dominant systematic is uncertainty in relating cluster masses
to observables. To tackle this obstacle, we are carrying out rather an empirical approach,which is to investigate nearby, massive clusters in great detail combining various techniques(optical, cluster lensing, X-ray and the Sunyaev-Zel’dovich eﬀect) in order to derive thewell-calibrated mass-observable relation (Miyazaki et al. 2002; Okabe et al. 2007).
Self-calibration
An alternative promising to protect the cosmological constraints against the systematics
is combining various dark energy probes, because diﬀerent methods depend on cosmologicalparameters in diﬀerent ways and also are sensitive to diﬀerent systematics. If properties ofthe systematics inherent in each probe are fairly well modeled, the combined measurements
can be a powerful way to constrain the cosmological parameters and the model systematics
simultaneously, reaching to the self-calibration regime. For the forecasts above, we focusedon two observables drawn from the HSC WL survey: weak lensing tomography and cluster
counts. There are various other dark energy probes available from the same HSC WL survey:
the higher-order correlation functions of the cluster distribution and the cosmic shear ﬁeld
(e.g. the lensing bispectrum studied in Takada & Jain 2004), the information inherent in
the higher-order distortion eﬀect on galaxies images beyond the shear (Okura et al. 2007),and the baryonic acoustic oscillation experiment of angular galaxy power spectra using thephoto-z information (Dolney et al. 2006). It will be worth exploring an optimal combinationof various probes such that the combination can maximize the dark energy constraints aswell as protect the constraints against the systematics in a most secure way.

2.3.1.3

A prospect for testing the gravity theory

The gravity is the fundamental force that governs the nature of the universe at distantscales. Many experimental tests in the solar system have been done, and it is known thatthe general relativity is the theory consistent with the experiments. Then, the discovery of
the accelerated expansion of the universe is quite interesting because it means the repulsive
force of the gravity at the cosmological scales. Thus, the accelerated expansion of the universe
might be originated from the modiﬁcation of the gravity theory itself, instead of the presenceof the dark energy component in the universe. Only the cosmological observations will giveus the clue to explore the origin of the accelerated expansion of the universe. At the sametime, the cosmological observations give us the new constraint on the gravity theory at thecosmological scales. The aim of this subsection is to investigate a prospect for testing thegravity theory with the future HSC WL survey.
The experimental test of the general relativity has a long history. The parameterised postNewtonian (PPN) framework played an important role in the experimental test, because
it gives a systematic description of the deviation from the general relativity. Then, thetheoretical framework for such the systematic description of the cosmological perturbationis considered, which is called the cosmological parameterised post-Newtonian framework or
the parametrised post-Friedman framework. In general, the metric of the 4 dimensional
spacetime with the scalar mode perturbation can be written as, in the conformal Newtonian2.3. SCIENCE GOALS OF HSC WL SURVEY

23

gauge,£¤
ds2 = a2 (η) − (1 + 2Ψ) dη 2 + + (1 + 2Φ) d~x2 ,where a(η) is the scale factor, η is the conformal time, Ψ and Φ are the gravitational potential
and the curvature perturbation, respectively. In the general relativity, in the case of noanisotropic stress and no entropy perturbation, we have,Φ + Ψ = 0,∆Φ = −4πGa2 ρm δ,
on sub-horizon scales, where G is the gravitational constant, ρm is the matter density, and δ
is the perturbation of the matter density contrast. In a modiﬁed gravity theory, the above
relation can be modiﬁed by introducing Γ and Geﬀ asΦ + ΓΨ = 0,∆Φ = −4πGeﬀ a2 ρm δ.
Here we investigate a future prospect of constraining the modiﬁed gravity model by adopting the cosmological parameterised post-Newtonian framework proposed by Bertschinger andZukin (2008). In the general relativity, it is known that the evolution of the long wavelength
cosmological perturbations are determined from the background Friedman equation and the
local energy-momentum conservation lows. Generalising this to arbitrary theories of gravity,
Bertschinger (2006) investigated the consistency relation of the cosmological perturbations,being valid for the 4-dimensional metric theory of gravity with a well-deﬁned infrared (longwavelength) limit and with the local energy-momentum conservation. Following Bertschinger
and Zukin, we adopt the parameterisation
Γ(a) = 1 + βas ,with the assumption of the scale-independent modiﬁcation, and β and s are the parametersof modiﬁcation. Because of the consistency relation, Γ and Geﬀ are related with each othereﬀectively, then Geﬀ is determined by Γ(a).
Figure 2.8(a) plots the evolution of the growth factor of the (modiﬁed) gravity model
speciﬁed by β and s. Figure 2.8(b) is the same as ﬁgure 2.8(a) , but the growth rate as
the function of the scale factor. The diﬀerence of the growth of the perturbation can betested with the cosmological observation of the weak lensing statistics and the combination
with a baryon acoustic oscillation (BAO) survey. Figure 2.9 exempliﬁes the constraint fromfuture surveys. The solid curve in Figure 2.9 shows the 1-sigma conﬁdence contour in theβ − s plane, which is obtained by marginalizing the Fisher matrix of the 8 parameters,β, s, Ωm , Ωb , h, ns , and σ8 over the parameters other than β and s. Here the target
parameters are β = 0.5, s = 1 (marked), Ωm = 0.28, Ωb = 0.044, w = −1, h = 0.7,
ns = 0.95, σ8 = 0.8. For the sample, we assumed the survey area of 2000 square degrees,with the number density 30 galaxies per square arcminutes, and the tomography of the 4
redshift bins. The dash-dotted line shows the general relativity (β = 0). The dashed curve isthe combination with an additional redshift survey like the WFMOS BAO survey. Here, forthe galaxy redshift survey, we assume that the same survey area, 2000 square degrees, andthe range of redshift 0.8 ≤ z ≤ 1.4, and the number density 4 × 10−4 h3 Mpc−3 . In modeling24CHAPTER 2. INSTRUMENT REQUIREMENTS FROM WEAK LENSING SCIENCES

Figure 2.8: (a, Left) Growth factor divided by the scale factor, D1 (a)/a, as function of thescale factor a (see also [3]). The solid curve is the general relativity, the dashed curve is
β = −0.5 and s = 1, and the dotted curve is β = +0.5 and s = 1. The other parameters areΩm = 0.28, ΩΛ = 0.72. (b, Right) Same as (a), but the growth rate, d ln D1 (a)/d ln a.

Figure 2.9: The 1-sigma conﬁdence contour in the β − s plane. The solid curve is theconstraint from the weak lensing and the dashed curve is the constraint from the combinationof the weak lensing and the BAO measurement assuming an additional redshift survey. Thissolid curve is the result with the 8 parameters of the Fisher matrix for the lensing power
spectrum and the dashed curve is the 1-sigma constraint combined with the galaxy power
spectrum from the redshift survey. The target model is β = 0.5 and s = 1.

2.3. SCIENCE GOALS OF HSC WL SURVEY

25

the galaxy power spectrum of the redshift survey, we simply considered the linear theory
adopting the scale independent bias model. Figure 2.9 demonstrates that the weak lensing
survey will give some useful constraint on the modiﬁed gravity, and that the combination
with the redshift survey improves the constraint signiﬁcantly. But, note that the constraint
depends on the sample, especially, the survey area and the photometric redshift error areimportant.
The test of the general relativity will be an important task of future large surveys. Possible
signature of modiﬁcation of gravity theory would appear in the cosmological perturbations,to which the weak lensing statistics can be sensitive. Then, the HSC WL survey will provide
us with a chance to test the gravity theory. The HSC WL survey alone might not give astringent constraint, but it is important to establish the strategy for testing gravity theories.

2.3.2

Weak lensing cluster survey

2.3.2.1

Expected weak lensing cluster counts

Let us start with reviewing the ability of searching for clusters of galaxies in a wide ﬁeld
survey via the weak lensing technique. In subsequent subsections, we will discuss sciencecases which are feasible with a cluster catalog constructed from the HSC WL survey.
Since the ﬁrst discovery of a shear selected cluster with the spectroscopic conﬁrmation
by Wittman et al. (2001), there has been a lot of progress in constructing a mass selectedcluster catalogue from unbiased survey using the weak lensing technique. The Suprime-Cam
has made a pioneering contribution to this ﬁeld. Miyazaki et al. (2003) ﬁrst reported thedetection of several signiﬁcant cluster candidates in an unbiased 2.1 deg2 ﬁeld. The ﬁrst
sizable samples of weak lensing shear selected cluster candidates were recently presented byMiyazaki et al. (2007). Their sample selection was based solely on high peak ﬁnding in
weak lensing density maps, and they found 100 signiﬁcant peaks in a 16.7 deg2 survey area.
They also tested the success rate by comparing their weak lensing cluster candidates in the
XMM-LSS ﬁeld (Figure 2.10) with X-ray clusters taken by XMM-Newton (Valtchanov et al.,
2004; Willis et al., 2005; and Pierre et al., 2006), and found that 12 out of 15 weak lensing
cluster candidates are also identiﬁed by X-ray survey. Thus the success rate estimated inthis way is 0.8 (12/15).
The detection rate obtained by the Suprime-Cam weak lensing survey was found to beconsistent with the theoretical prediction based on the currently standard ΛCDM model
made by Hamana, Takada & Yoshida (2004). They used cosmological N -body simulations
as well as the analytic models to investigate the ability of a weak lensing cluster survey, andfound that for the survey parameters of Miyazaki et al. (2002; 2007), weak lensing survey
can locate massive clusters of M > 1014 M¯ h−1 at 0.1 < z < 0.7 (Figure 2.11), and the
expected number of detectable clusters is 4-5 per one square degrees (Figure 2.12), which isindeed consistent with the observed value.
Since the (observationally and theoretically) expected cluster number counts is about 5
per one square degrees, the expected total number from the proposed 2000 deg2 HSC WL
survey is ten thousand including more than one hundred most massive (M > 1015 h−1 M ¯)
clusters, which will be the ﬁrst statistical mass selected cluster sample. Let us comment on
an advantage of the HSC WL survey over other shallower survey. As compared in Figure
2.11 the HSC WL survey can locate about four times less massive clusters than the shallower
surveys, and, as its consequence, the number density of the HSC WL survey is about one

26CHAPTER 2. INSTRUMENT REQUIREMENTS FROM WEAK LENSING SCIENCES

1

9

14
4

20

8
11
18
13

0

8

17

22 10 12

23

21
5

3
15

6
16 19

2
7

Figure 2.10: The κ-S/N map for the XMM-wide ﬁeld (thick red contour). The lowest contour
is set at S/N = 2 with the increments of 0.5. The smoothing scale, θG = 1 arcmin. The thinblue contours show the surface number density of visible (21< Rc <23) galaxies. Positionsof detected halos (S/N ≥ 3.7) are designated by small circles. Solid circles represent secure
halos (see Miyazaki et al. 2007) whereas unﬁlled circles represent samples detected outside
the secure area. Squares show the location of X-ray selected clusters of galaxies publishedin Valtchanov et al. (2004), Willis et al. (2005) and Pierre et al. (2006). Large open circles
show halos followed up by FOCAS and LRIS (see Miyazaki et al. 2007 for details). The RC band imaging data, from which the κ map was made, consists of 13 Suprim-Cam pointings
illustrated (very roughly) by the green rectangles. The orange circle show the expected HSC
ﬁeld of view of 1.5 degree φ.

2.3. SCIENCE GOALS OF HSC WL SURVEY

27

Figure 2.11: ΛCDM prediction of the region in a plane of the cluster redshift and cluster
mass where a cluster of galaxies is expected to be located by the weak lensing technique .
The red region is for HSC WL survey (ng = 35 arcmin −2, hzs i = 1), whereas the green
region is for a shallower survey (ng = 13 arcmin−2 , hzs i = 0.7).

Figure 2.12: ΛCDM prediction of the mass distribution of clusters of galaxies located by
the weak lensing technique . The red curve is for HSC WL survey (ng = 35 arcmin −2,
hzs i = 1), whereas the green curve is for a shallower survey (ng = 13 arcmin−2 , hzs i = 0.7).

28CHAPTER 2. INSTRUMENT REQUIREMENTS FROM WEAK LENSING SCIENCES

Figure 2.13: The bold solid contours show the projected mass density ﬁeld reconstructed fromthe weak lensing measurement based on the Suprime-Cam data of the famous massive cluster
A1689, superposed onto the Subaru image. The thin solid contours show the distribution ofmember galaxies selected from the red sequence in the color and magnitude diagram. See
Umetsu, Takada & Broadhurst 2007 for the details.

order of magnitude larger than a shallower survey (Figure 2.12). This is due to the highgalaxy number density (ng = 35 arcmin−2 for HSC WL survey, whereas ng = 13 arcmin−2for shallower survey) and a higher mean redshift of galaxies. Thus in the next 10 years, only
the HSC WL survey will provide a sizable (N ∼ 105 ) catalog of mass selected typical massclusters.
For a visual impression of how eﬃcient the HSC weak lensing cluster survey will be, we
compare in Figure 2.10 the ﬁeld-of-views of the Suprime-Cam (the green rectangles) with
one expected for the HSC (assuming 1.5 degree φ shown by the orange circle). The imaging
data of the XMM-LSS ﬁeld is made from 13 Suprime-Cam pointing as shown in the plot andits eﬀective area is 2.8 deg2 . The HSC will be able to cover the same area (actually evenslightly larger area) by only two pointings. In the plot, one may ﬁnd that many clusters
are located in the one HSC pointing. This illustration may indeed impress us that one may
locate about 5-10 clusters from a single HSC pointing of ∼20 minutes exposure.
2.3.2.2

A lensing study of cluster mass distribution

A natural consequence of small-scale, non-linear gravitational clustering in the CDM structure formation scenario is appearance of dark matter halos that are hosting galaxies andclusters of galaxies, most important objects in astronomy. In particular clusters of galaxies

2.3. SCIENCE GOALS OF HSC WL SURVEY

29

are the most massive, gravitationally bound objects in the Universe, and therefore gravity mainly due to dark matter is playing an essential role to their formation and evolution.
While an N-body simulation provides a most powerful means for investigating the non-linear
structure formation, the most intriguing prediction is appearance of characteristic mass distribution in a dark matter halo: the mass density proﬁle of simulated halos can be well ﬁtted
by one simple function over a wide range of mass scales from galactic to cluster scales. This
universal density proﬁle is called Navarro-Frenk-White model (NFW; Navarro, Frenk andWhite 1996, 1997) given by ρ(r) ∝ r−1 (1 + r/rs )−3 with logarithmic slopes ρ ∝ r−1 and r−3
at radii smaller and larger than the scale radius rs . In addition the shape of dark matter
halos is found to be triaxial rather than spherical even in a statistical average sense (Jing &Suto 2002), originating from gravitational tidal interaction by the surrounding ﬁlamentary
cosmic webs. These properties of dark matter halos should be sensitive to the nature ofdark matter such as collision-less nature and negligible thermal velocity dispersion on cosmological scales, at least scales relevant for clusters, which are usually a priori assumed insimulation-based studies. Therefore an observational test of the dark matter halo properties
is extremely important in order to derive, albeit indirect, constraints on the nature of dark
matter, independently of laboratory experiments.
Gravitational lensing oﬀers a unique way for recovering the total matter distribution in acluster, irrespective of bias uncertainties arising when dark matter distribution is inferredfrom visible baryonic matter distribution (galaxies and intracluster gas). For a cluster region,
the gravitational lensing shearing eﬀects on the images of background galaxies are relatively
easily measurable compared to cosmic shear, allowing a reconstruction of the mass distribution up to scales of the virial radius. However, in implementing these lensing studies, it iscritically important to make a secure selection of background galaxies because only galaxies behind a given cluster are gravitationally lensed, as clearly pointed out in Broadhurst,
Takada, Umetsu et al. (2005) (also see Medezinski, Broadhurst, Umetsu et al. 2007; Okabe,
Takada, Umetsu et al. 2008). Otherwise, an inclusion of member galaxies or foregroundgalaxies may cause a signiﬁcant dilution of the measured lensing signals11 . A multi-color
data set of the HSC survey will be deﬁnitely useful to accurately select background galaxies
as well as estimate average redshift of the background galaxies, which are in turn importantto estimate, in an unbiased way, the physical quantities of cluster such as the viral mass. We
will below describe in more detail science cases for cluster-related physics that are possible
to carry out using the HSC WL survey.
Lensing study of individual massive clusters
As stated in detail in § 2.3.2.1, the HSC survey enables us to, based on the weak lensing
measurements, construct a mass-selected cluster catalog that is expected to contain ∼ 105>
1014 h−1 M¯ . In particular, for massive clusters in the catalog,
halos with virial masses M ∼the depth of the HSC survey or more exactly the number of background galaxies available forthe weak lensing analysis, allows us to make a detailed lensing study of individual clusters.
In particular, it would be interesting to deﬁne sub-samples from the cluster catalogs basedon the properties of the mass distribution: a straightforward one is to deﬁne sub-samplesof ‘relaxed’ and ‘unrelaxed’ clusters depending on whether or not the reconstructed mass
11

In particular the dilution eﬀect is more signiﬁcant in the inner region of cluster as member galaxies areusually more centrally concentrated than the dark matter distribution, thereby causing a bending of theradial shear proﬁle at larger radii than it actually is. For example, this dilution causes an underestimationin constraining the concentration parameter from the measured shear proﬁle.

30CHAPTER 2. INSTRUMENT REQUIREMENTS FROM WEAK LENSING SCIENCES

SubaruSubaru+Hubble

Figure 2.14: Left panel: The constraint errors on the concentration parameter cvir and thevirial mass Mvir derived by comparing the weak lensing measurements for A1689 with the
NFW predictions. The cross mark denotes the best-ﬁt model, while the three contours
correspond to ∆χ2 = 2.3, 6.17, 11.8, respectively. The dashed curve indicates the NFW
models that reproduce the observed Einstein radius of θE ≈ 4000 for zs = 1. Right panel: It
is demonstrated that adding the strong lensing constraints into the weak lensing constraints
can signiﬁcantly improve the parameter constraints. This plot is taken from Umetsu &Broadhurst (2007).

distribution displays a major merger phenomena or many sub-clumps. The feasibility ofthese studies have been already proven by using the Suprime-Cam data sets, and the HSC
survey can further signiﬁcantly improve the statistical precision of these studies, as will be
described in detail below along with the actual examples.
A sample of relaxed clusters would be more appropriate to quantitatively test the NFW
proﬁle or the triaxial halo model, because the theoretical predictions have been usually derived using simulated halos that look well-relaxed halos. Figure 2.14 shows such an example,
showing the mass distribution reconstructed of the famous massive cluster A1689, derived
from the weak lensing measurements of the Suprime-Cam i0 -band data. The Subaru deep
image thus allows a signiﬁcant detection of the weak lensing signals, more than 10σ for the
total signal-to-noise ratio in this case. It should be worthwhile noting here that, to obtain
the well-isolated mass distribution like that shown here, we need to measure the shearing
eﬀects on background galaxy images up to large radii, typically the virial radii (a few Mpc
or a few 10 arcminutes on the sky), because the weak shear is a non-local quantity, therefore
non-vanishing even in the region where the mass density is locally zero, and the signal is only
slowly decreasing with radius from the cluster center. Therefore, having a large ﬁeld-of-view2.3. SCIENCE GOALS OF HSC WL SURVEY

31

is extremely important for a detailed study of cluster lensing.
Once secure a lensing measurement is obtained, we can directly estimate the virial massby comparing the model predictions of the mass distribution with the measured signal. Figure 2.14 shows such an example for A1689. In this case, the measured lensing signals arewell ﬁtted by an NFW proﬁle which is speciﬁed by two free parameters, e.g. the virial mass
Mvir and the concentration parameter cvir (Broadhurst, Takada, Umetsu et al. 2005; Umetsu& Broadhurst 2007). How often an NFW proﬁle can give a reasonably good ﬁt to actual
clusters is still not so clear, so this is the issue that can be quantitatively addressed with
the HSC survey. As shown in the left panel of Figure 2.14, the weak lensing measurement
alone cannot well constrain both cvir and Mvir generally leaving degeneracies between thetwo parameters. A powerful way to overcome this limitation is to add strong lensing information that puts stringent constraint on the mass distribution in the innermost region.
Observationally, we have a very good chance to ﬁnd strong lensing events in the inner region
from the deep imaging data, especially if the cluster is massive enough. As demonstrated
in the right panel, combining weak and strong lensing can eﬃciently break the degeneracy,
leading to stringent constraints on Mvir and cvir separately.
As can be seen in Figure 2.13, the cluster mass distribution is not circularly symmetric. It
would be interesting to explore whether the observed mass distribution can be better ﬁtted
by the triaxial halo model than by the spherical NFW proﬁle (Oguri, Takada, Umetsu et al.
2006). This would give additional constraints on the nature of dark matter. Also quantifying
deviations of individual cluster mass distribution from the average properties would be aninteresting direction to explore.
Also it would be interesting to study how the total matter distribution is related to thespatial distribution of member galaxies within a cluster, i.e. the galaxy bias. This study canbe done as a function of galaxy types, which would be useful to develop our understanding of
the assembly histories of diﬀerent types of galaxies within a cluster. Similarly, if the X-ray or
SZ observations are available, it is also interesting to explore the relations between the total
matter, intracluster hot gas and galaxy distributions, in order to obtain a comprehensive
picture of cluster formation and evolution processes.
Alternative interesting direction to be explored on individual galaxy basis is to make adetailed study of ‘merging’ clusters combining all the observables, the weak lensing, galaxy
distribution, X-ray and SZ observations. The dark matter particles are believed to becollision-less, we know galaxies are collision-less due to the diﬀuse number density, while
the intracluster hot gas is collision causing shocks in the merging stages. A very interesting
example has been recently provided from the detailed lensing studies of the so-called Bullet
Cluster, IE 0657-56 (Clowe et al. 2006). They found signiﬁcant oﬀsets between the distributions of the collision-less components (lensing matter and galaxies) and the intracluster
hot gas, providing an indirect evidence of the existence of the collision-less dark matter.
Such merging clusters where major mergers are going on in the plane perpendicular to theline-of-sight on the sky would be very rare, e.g. may be like a one event of 100 cluster-sizedhalos, without the projection eﬀect, as estimated from simulation-based studies (Hayashi &White 2006). The HSC survey will deﬁnitely increase such sample of merging clusters. In
Figure 2.15, Okabe & Umetsu (2008) showed that A520 may be another example of bulletlike clusters because the oﬀset between the collision-less components and the hot gas arefound, at less signiﬁcance. Thus, investigating merging clusters in great detail with weaklensing would be a very interesting topic.

32CHAPTER 2. INSTRUMENT REQUIREMENTS FROM WEAK LENSING SCIENCES

Figure 2.15: Upper-left panel: The Subaru i0 -band image of A520, superposed with themass density contours derived from the weak lensing measurements. The plotted region is≈ 120 × 120 . Upper-right: The contours show the adaptively-smoothed Chandra X-ray (0.77keV) image, overlaid on the same i0 -band image. Lower-left: The luminosity density ﬁeldof member galaxies, smoothed to the same angular resolution to that of the lensing mass
map. Lower-right: The X-ray image overlaid on the mass density contours. This plot istaken from Okabe & Umetsu (2007).

A statistical study of cluster mass distribution: stacked cluster lensing
A novel approach to cluster lensing which has recently got much attention is the so-called“stacked” cluster lensing: the average cluster lensing signal is obtained by stacking thetangential shearing eﬀect on background galaxies, with respect to cluster center, over a large
number of clusters selected with observational properties (optical richness, individual cluster
lensing signal, X-ray brightness and so on). The HSC weak lensing survey may allow us>
1014 h−1 M¯ based on their optical
to identify ∼ 105 cluster candidates with masses M ∼richness and individual lensing signals as described in § 2.3.2.1. A notable advantage ofstacking is that, as long as the Universe is homogeneous and isotropic, a very smooth shear
proﬁle can be obtained, allowing a cleaner comparison between the measurement and thetheoretical prediction, the spherical NFW proﬁle, while the asphericity and substructures ofcluster mass distribution and the projection eﬀects otherwise give severe contamination to
the comparison on individual cluster basis. Moreover, in contrast to SDSS and a shallower

2.3. SCIENCE GOALS OF HSC WL SURVEY

33

Figure 2.16: The expected “stacked” tangential shear proﬁle for clusters obtained whenclusters having the individual cluster signals greater than (S/N ) ≥ 5, selected for the HSC
weak lensing survey selected from the HSC weak lensing survey. The error bars show 1-σstatistical errors due to the sample variance, contamination due to uncorrelated large-scale
structures along the line-of-sight.

34CHAPTER 2. INSTRUMENT REQUIREMENTS FROM WEAK LENSING SCIENCES

Figure 2.17: AMiBA SZE and Subaru/S-Cam observations of the cluster A1689 at z = 0.183:(Left) SZE sky map imaged with the 7-element AMiBA. (Right) Lensing convergence mapwith the spin-2 gravitational shear ﬁeld overlayed. In both panels, the ﬁeld size is taken to
be 220 for comparison. For each panel the Gaussian FWHM in map-making is indicated by
the solid circle at the bottom right corner.
survey, the HSC survey of the proposed depth allows us to quantify the mean and scatters in
mass-observable relations, since the cluster mass can be estimated on individual basis fromthe lensing signal.
Fig. 2.16 shows the expected cluster lensing signal obtained when cluster regions having
the lensing signals greater than (S/N ) ≥ 5 are stacked. Clearly one can expect signiﬁcantsignal-to-noise ratios for the stacked lensing signal measurement; the errors at each radialbins can be reduced by a factor of 10 compared to the individual cluster lensing. The
two thin curves show the NFW proﬁle contribution of individual cluster mass distribution
and the contribution of the ﬁlamentary structure surrounding. The latter contribution issensitive to halo bias parameter that can be accurately modeled in terms a given cosmologicalmodel, i.e. cosmological parameters (e.g., Mo & White 1996). The structure formation
model predicts that, even if the same redshift of lenses is considered, more massive halos are
more biased objected having greater halo biases, therefore the greater shear contributions
from the surrounding ﬁlamentary structures are expected. Thus extracting the halo bias
information may give additional, clean cosmological information in addition to the cluster
counting statistics. Also intriguing that was recently proposed by Dalal et al. (2007) is thatthe halo bias may be another vital way to probe the primordial non-Gaussianity. Using the
ray-tracing simulations, we will carefully investigate the feasibility of the HSC weak lensing
survey for testing the NFW proﬁle as well as for constraining cosmological parameters andthe primordial non-Gaussianity from the halo bias measurements.

2.3.2.3

Synergy with SZE Observations

In recent years a lot of eﬀort has been put into developing sensitive radio/submm telescopes
dedicated to observations of the CMB temperature and polarization anisotropies and theSunyaev-Zel’dovich eﬀects (hereafter SZE) for cosmology and structure formation. Princeton

2.3. SCIENCE GOALS OF HSC WL SURVEY

35

University is a major partner of the ongoing ACT (Atacama Cosmology Telescope) project,12and ASIAA has been leading the AMiBA (Array for Microwave Background Anisotropy)
project.13
Large radio interferometers, such as BIMA, OVRO, and NRO, have been traditionallyused for the measurements of the thermal SZE towards known clusters of galaxies (e.g.,
Birkinshaw 1999; Carlstrom, Holder, & Rees 2002). Although these radio interferometers
are sensitive to small structures, their ﬁeld of view (FoV) is rather small (say, a few arcminutes) and hence not optimized for a large scale survey of clusters and large scale structure.
The new generation CMB/SZE radio interferometers, such as AMiBA at 94GHz (Arrayfor Microwave Background Anisotropy), CBI at 30GHz (Cosmic Background Imager), and
SZE at 30/94Ghz (Sunyaev-Zel’dovich Array), are speciﬁcally designed to study large scale
structures on the CMB background with large FoVs (say, 10-20 arcminutes) and the angularresolutions optimized for the detection of high-z clusters (a few arcminutes). The AMiBA
is a 86–102 GHz interferometer array sited at Mauna-Loa, Hawaii, speciﬁcally designed tosample structures in CMB as small as 20 in size (Ho et al. 2004). In 2007 AMiBA observed 6
Abell clusters via the thermal SZE with seven 0.6m diameter antennas in a compact conﬁguration. The expansion for thirteen 1.2m antennas is on the way to improve the sensitivity
and angular resolution. As demonstrated by DASI, CBI, and OVRO/BIMA observations,the inherent power of interferometers is to be able to reject uncorrelated noise and control
systematic, allowing fairly deep imaging on a targeted ﬁeld. AMiBA will be therefore usefulfor follow up SZE observations of cluster candidates detected by the HSC WL survey. Figure
2.17 compares the AMiBA SZE and Subaru WL sky maps derived for the cluster A1689 at
z = 0.183.
In the past several years, on the other hand, increasing attention has been given to thedevelopment of SZE-dedicated, multi-pixel, multi-band bolometer arrays, such as ACT (Atacama Cosmology Telescope), SPT (South Pole Telescope), and APEX (Atacama Pathﬁnder
Experiment). These new generation bolometer arrays are not only highly sensitive but also
have a very wide FoV of ∼ 1deg, optimized for large cluster surveys. The ACT is a 6m
millimeter-wave telescope (145, 220, 265 GHz) ﬁelded with the Millimeter Bolometric Array Camera. ACT is located at an altitude of 5200 meters in the Chilean Andes, and isdesigned and dedicated to conduct large sky surveys with high sensitivity to the SZE andCMB temperature and polarization anisotropies. The ACT SZE survey is expected to cover400deg2 in the southern sky, with 2µK sensitivity per 20 pixel. In contrast to interferometers,
the power of multi-pixel bolometric arrays is their wide-ﬁeld imaging capability with superb
sensitivities, allowing to conduct large sky surveys, providing a statistically large sample ofSZ-ﬂux limited clusters.
Both blind SZE and WL surveys aim to detect massive clusters of M > 1014 M¯ /h tostudy the formation of structure and test the cosmological model, through the abundance
and properties of the detected clusters. Since the cluster SZE ﬂux is a measure of the total
thermal energy (integrated thermal pressure) in the intra-cluster medium (hereafter ICM),
it is rather insensitive to the cluster core as compared with the X-ray measurement, andhence provides a well-understood mass-observable relation. Furthermore the SZE brightness
is redshift independent. This nature of the SZE makes SZE surveys essentially mass-limitedwith little redshift dependence, as demonstrated in Figure 2.18 (Umetsu et al. 2004). A
combination of WL and SZE surveys will not only help to improve the cluster detection
1213

http://www.physics.princeton.edu/act/
http://amiba.asiaa.sinica.edu.tw/

36CHAPTER 2. INSTRUMENT REQUIREMENTS FROM WEAK LENSING SCIENCES

Limiting mass, Mhalo [MO• h-1]

1015
SZE peak flux limit
3σ
4σ
4.2σ

10% false rate at 4.2σ
1σ = 0.8mJy (94GHz)

1014
0.0

0.2

0.4
0.6
Redshift, zhalo

0.8

1.0

Figure 2.18: Example of the SZE limiting mass as a function of halo redshift derived foran AMiBA deep SZE survey based on preheated N -body/hydrodynamic simulations in theΛCDM model, normalized to reproduce the X-ray scaling relations in the local universe.
Shown are the predictions for three diﬀerent SZE peak ﬂux limits: 3σ (dashed), 4σ (dotted),and 4.2σ (solid), with the rms sensitivity of 1σ = 0.8 mJy at 94GHz. At 4.2σ, the false
positive rate is 10%.
eﬃciency by utilizing the cross correlation between the WL and SZE signals, but also reveal
much about the diﬀerent structures of mass (dark matter) and hot baryons in clusters of
diﬀerent mass (see §2.3.2.4)
2.3.2.4

Synergy with X-ray observations

Most of baryons in clusters are in the form of the hot-thin plasma with temperature ∼a few − 10 keV and the density ∼ 10−3 − 10−2 cm−3 , called the intra-cluster medium (ICM).
The ICM accounts for ∼ 10−15 percent of total mass and thus the non-gravitational process,such as radiative cooling, and the hydrodynamics in the ICM play a key role in the formationand evolution of clusters. Our understating of clusters has been provided for a long time by
X-ray satellites, ROSAT, ASCA, XMM-Newton, Chandra and Suzaku. Especially, Chandra
and XMM-Newton enable us to measure the density, temperature, pressure, entropy andmetal abundance of the ICM with high resolution. However, we have not yet understoodthose ICM phenomena in detail, because the distributions of dark matter which controls
the dynamics of the ICM have not been revealed directly by X-ray studies. Therefore,a systematic joint weak-lensing and X-ray analysis is a unique method to investigate thephysical interplay between the ICM and dark matter, and establish mass-observable relationsof clusters.
The WL survey unveils mass clumps in the large-scale structure and the substructuresin massive clusters (M > 1014 h−1 M¯ ) as demonstrated in Hamana et al. (2008), Okabe &2.3. SCIENCE GOALS OF HSC WL SURVEY

37

Umetsu (2008). Reconstructed mass maps provide a new approach on X-ray analysis : (1) to
set an accurate pointing and an adequate exposure time in order to observe eﬃciently X-ray
emitting, and (2) to extract a spectrum in the region associated with mass. In practicalwork, we can investigate how much the ICM parameters (e.g. temperature and pressure)diﬀer from within/outside of mass structures. Thus, the spectral observations in segmentsof mass substructures would interpret better the origin of the local properties of the ICM.
And hence a joint analysis with WL study enables us to investigate the correlation between
the ICM parameters and mass. The ﬁlamentary structures or galaxy groups embedded in
the large-scale structure, revealed by the WL analysis, is a clue to detect, for the ﬁrst time,
the warm-hot intergalactic medium (WHIM) thanks to both accurate pointing and eﬃcient
spectral analysis. The WHIM with temperatures 105 − 107 K is supposed to distribute inthose structures and contain a signiﬁcant fraction of baryons. It might become the powerful
technique to search a missing baryon.
During the operation of HSC survey, the ﬁve X-ray satellites, Chandra, XMM-Newton,
Suzaku, Astro-H (NeXT) and Spektrum-Roentgen-Gamma (SRG), would be in orbit andobserve the diﬀuse X-ray emitting gas. The high angular resolutions of Chandra and XMMNewton is powerful for measuring the ICM parameters within r2500 − r500 for clusters, although we cannot rule out a possibility that they become superannuated. The Japanesesatellite, Suzaku, with the low and stable X-ray background but low angular resolution(FWHM∼ 20 ), is the unique tools to detect the low X-ray ﬂux from groups or outskirts<<(r500 ∼r∼rvir ) of clusters, because observations of those regions does not in principle require the high angular resolution. The Suzaku with WL study would lighten the world of
low X-ray ﬂux. It notes that the FoV of HSC could cover the viral radius of nearby massive<
cluster (z ∼0.1) with one pointing and therefore determine the viral mass without the assumption of dynamical state. The X-ray satellite, on the other hand, has the ﬁducial range of<<a radius for studying the ICM (r ∼r2500 − r500 for Chandra and XMM-Netwton and r500 ∼rfor Suzaku). We therefore select the best satellite for scientiﬁc goal or combine them in order
to compare WL results within a virial radius. The Russian satellite, Spektrum-RoentgenGamma (with the imaging instrument eROSITA), designed for all sky survey would launch
in 2011. The SRG would systematically detect diﬀuse X-ray emissions from mass structures,
which is complementary to the WL survey. The main scientiﬁc goal of the next Japanese
X-ray mission, Astro-H (old name is NeXT), which will be launched in 2013, is to detect
the turbulence in the ICM and map the hard X-ray emission out. They are closely relatedwith cluster mergers, and hence the mass distributions are complementary constrains on thephysical properties of the ICM. We could also conduct the same science as Suzaku with theimaging instrument of Astro-H.
2.3.2.5

Cross correlating cosmic shear with large scale structure for probing
the galaxy/cluster bias

The relation between the distributions of visible galaxies and underlying mass is one of the
key issues for understanding the formation of large scale structure and the galaxy formation
scenario. Observational work on the clustering statistics from galaxy surveys has shownthat clustering properties are diﬀerent among galaxies, quasars, and galaxy clusters, andthat galaxies of diﬀerent morphologies (elliptical, spiral), luminosities, and spectral (red,blue) types have diﬀerent clustering properties. This leads to the implication that galaxies
are not a perfect tracer of mass, or the galaxy distribution is biased against the underlying

38CHAPTER 2. INSTRUMENT REQUIREMENTS FROM WEAK LENSING SCIENCES

6050
HSC WL Survey
DES

50

HSC

r / σ(r)40

35

HSC WL Survey
DES30

b=1, r=1

30

40

2

r / σ(r)Survey area = 2000 deg
2
nb = 35/arcmin.
2
nf=20/arcmin.
zs=1.0HSC WL Survey
2Survey area = 2000 deg
2
nb = 35/arcmin.
2
nf=20/arcmin.
zs=1.0

45

25
20

20

10

0

15

DES
2
Survey area = 5000 deg
2
nb = 12/arcmin.
2
nf=8/arcmin.
zs=0.681

10

θap[arcmin.]DES
2
Survey area = 5000 deg
2
nb=12/arcmin.
2
nf=8/arcmin.
zs=0.68

10
5
100

0

0

0.2

0.4

zf

0.6

0.8

Figure 2.19: Expected performance of HSC WL survey to explore the relation between
mass and galaxy distributions. (Left) Shown are HSC sensitivities r/σ(r) for the correlationcoeﬃcient between galaxy and mass, r = hδm δg i/(σm σg ), as a function of angular scale θap .
Here the model with b ≡ σg /σm = 1, r = 1 is adopted for the estimation. The solid anddashed curves represent the predictions of HSC WL survey and DES, respectively. The HSC
WL survey has an advantage at smaller angular scales as the results of the higher number
density and higher mean redshift of the source galaxies. At the larger angular scale, theconstraints become comparable or less due to the smaller survey area than DES. (Right)
Same as the left panel, but shown here is the sensitivity as a function of foreground redshift
calculated at a ﬁxed angular scale, θap = 100 , along with the prediction for DES. It is clearlyshown that HSC has superior sensitivity to higher lens redshifts thanks to the large photon
collecting area of the 8m Subaru telescope.
mass distribution.
On the other hand, a large number of theoretical investigations have been done to model
the relation between the galaxy and mass distributions (e.g. Mo & White 1996; Tegmark &Peebles 1998; Taruya & Suto 2000; Yoshikawa et al. 2001), revealing complex properties ofthis relation, such as scale dependence, time evolution, dependence on halo mass, formation
epoch, nonlinearity and stochasticity.
Although comparisons between theoretical models and observations have been done using
the galaxy surveys which provide only the information on the visible galaxy distribution,
weak lensing surveys provide more suitable samples to explore this problem since weaklensing of background galaxies is a direct probe of the matter ﬂuctuations.
Schneider (1998) proposed a method to constrain the statistical relation between galaxy
and mass distributions by cross correlating the aperture mass statistics from the cosmic shear
measurements and the observedqgalaxy distribution. This method gives a direct measure
2 i where δof (1) the bias parameter, b = hδg2 i/hδmm and δg are the mass and the galaxy
overdensity, respectively,
q and (2) the correlation coeﬃcient between galaxy and mass distri-

2 i, which quantiﬁes how much galaxy and mass distributions
butions, r = hδg δm i/ hδg2 ihδm
diﬀer.
Recently, this method has been explored and applied to existing large galaxy surveys by
several groups, such as RCS/VIRMOS-DESCART surveys (Hoekstra et al. 2003) and GaBoDS (Simon et al. 2007). However, the current observational constraints on the theoretical2.3. SCIENCE GOALS OF HSC WL SURVEY

20

39

30
HSC WL Survey
DES

HSC
2
Survey area = 2000 deg
2
nb = 35/arcmin.
zs=1.015

HSC WL Survey
DES

DES
Survey area = 5000 deg
2
nb = 12/arcmin.
zs=0.68r / σ(r)r / σ(r)

20

10
DES
Survey area = 5000 deg
2
nb = 12/arcmin.
zs=0.68

5HSC
2Survey area = 2000 deg
2
nb = 35/arcmin.
zs=1.0

2

2
14

10

b=3, r=0.8, Mass>10 solar mass

14

Mass > 10 solar massb=3, r=0.8

0

1

10

θap[arcmin.]

100

0

0

0.2

0.4

zf

0.6

0.8

Figure 2.20: Same as Figure 2.19, but for cluster sample. (Left) Shown are HSC sensitivities
r/σ(r) for the correlation coeﬃcients of mass and clusters, r = hδm δcl i/(σm σcl ), as a functionof angular scale θap . For the foreground cluster sample, the mass limit is set to Mlim =1014 M¯ /h, as expected for the planned SZE cluster surveys by e.g. ACT. Here the PressSchechter mass function is used to calculate the abundance of clusters, and b = 3, r = 0.8
are assumed as the cluster bias model. The dashed curve represents a prediction for theplanned 5000deg2 DES for comparison. (Right) Same as the left panel, but shown here is
the sensitivity as a function of foreground redshift calculated at θap = 400 , along with theprediction for DES. Again it is clearly shown that HSC has superior sensitivity to higher
lens redshifts thanks to the large photon collecting area of the 8m Subaru telescope.
models are rather weak, limited by the cosmic variance. Ongoing and upcoming large weaklensing surveys are expected to improve the statistical accuracy of the measurements of thebias parameters, advancing our understanding of this unresolved problem. The HSC WL
survey will provide a powerful dataset for this purpose.
Here we address the feasibility of constraining the statistical relation between galaxy andmass distributions from the HSC WL survey. Figure 2.19 shows forecasts of the HSC constraints on the correlation coeﬃcient r for galaxies. We assume a 2000 deg2 total survey
area and the mean number density of galaxies used for the cosmic shear measurements,
nb = 35/arcmin.2 with the mean redshift zs = 1.0.
In the left panel, the expected constraints (the coeﬃcients r divided by the expected error
computed based on the formalism by Fan (2003)) for diﬀerent exposure times are plotted
as a function of aperture radius θap . For comparison, the same quantity for the DES isshown with the dashed curve. Deep multicolor photometry of the HSC survey will allow us
to measure photometric redshifts for individual galaxies to a good accuracy (σz . 0.15, see§2.4). With this accurate photo-z information, we will be able to subdivide the foregroundgalaxy sample into subsamples with diﬀerent redshifts, with an interval of ∆z ∼ 0.3. In
the right panel, we show the predicted HSC sensitivity for the coeﬃcient r as a function ofredshift for the foreground subsamples, obtained for the aperture radius of θap = 100 .
From Figure 2.19 we ﬁnd that the constraints on r from the entire sample with a few %accuracy are expected at the optimal angular scale of θap ' 50 from the HSC WL survey.
Comparing with the DES, the HSC WL survey has an advantage at smaller angular scales
as the results of the higher number density and higher mean redshift of the source galaxies.
At larger angular scales, the constraints become comparable to or less than DES due to the40CHAPTER 2. INSTRUMENT REQUIREMENTS FROM WEAK LENSING SCIENCES
smaller survey area.
For individual redshift subsamples, 3 to 5% accuracies are expected for redshifts lower
than z ∼ 0.8. The HSC WL survey especially will have an advantage at higher redshifts
than DES, owing to the large photon collecting area of the 8.2m Subaru telescope.
Figure 2.20 is the same as Figure 2.19 but for a sample of mass-selected galaxy clusters.
We assume a mass limit of Mlim = 1 × 1014 M¯ /h as expected for the planned SZE cluster
surveys by e.g., ACT. Although the constraints are worse than galaxies due to the shot noise,
5 to 20% constraints are still expected for clusters. Similarly to the case for galaxies, the
HSC WL survey has a great advantage at higher redshifts as compared to DES.
The higher order statistics for the cross correlation will also be useful to explore moredetailed properties of the complex relation between mass and galaxies/clusters distributions(Schneider & Watts 2005). The feasibility study of this statistics is in progress.

2.4

Photometric redshifts

Photometric redshift (photo-z) is a technique which enables us to measure redshifts of distant galaxies only from photometric data. Now in the era of multi-passband surveys by
8-10m telescopes, it has become an essential tool in the key sciences including observationalcosmology and galaxy formation and evolution. In fact, the feasibility of those sciences
strongly relies on photo-z accuracies. Thus we should thoroughly discuss how we establish
our method to obtain photo-z to accuracy levels necessary to each of the science cases in thecourse of survey designing for HSC.

2.4.1

Requirements for HSC WL Survey

For our key sciences, the weak lensing (WL) survey, photo-z will be used to perform astudy of WL tomography. The characteristics of errors associated with the photo-z haveto be well understood to derive reliable results from the WL tomography analysis. The
photo-z error can be roughly divided into two categories. One is the statistical dispersion ofphoto-z (rms error) with respect to real redshifts. The other is so-called catastrophic error,a large systematic error which is not represented by Gaussian-like dispersion error. The
catastrophic error is mainly due to misidentiﬁcation between spectral features, e.g., 4000
Åbreak and Lyman break.
For the WL study, the catastrophic errors is much more sensitive than the statisticalerror of photo-z to resulting dark energy constraints. An oﬀset of 1% in the mean redshiftdetermined by photo-z corresponds to an approximately 100% error in the equation of stateof the dark energy parameter, w. On the other hand, in order to perform the WL tomography
by dividing the galaxy sample at intermediate redshift into 2 to 3 subsamples, we are stillrequired to achieve the statistical dispersion of order 0.1 (see §2.3.1).

2.4.2

Simulations of photo-z Based on Our Method

We would like to present results from our photo-z simulation to see the dependence of photo-z
accuracies on the combination of passbands. Our method employs a typical template ﬁtting
which searches for the best-ﬁt SED and the redshift of a galaxy by comparing the observed
SED of a galaxy with template SED’s by means of χ2 minimization (Furusawa et al. 2000).

2.4. PHOTOMETRIC REDSHIFTS

41

Figure 2.21: Results of the photo-z simulation using a template-ﬁtting method based on
the 4 passbands g 0 , r0 , i0 , z 0 . The relationship between input redshifts and output photo-z
for ∼ 2200 SEDs is shown in each panel. Diﬀerent symbols denote diﬀerent spectral types(open circles: disk-like, dots: spiral-like, and triangles: elliptical-like). The left panel shows
the result assuming photometric error of S/N=10 in each of the 4 passbands, while the rightpanel shows the result with S/N=30.

Figure 2.22: Same as in Fig. 2.21, but for the 5 passbands of g 0 , r0 , i0 , z 0 , Y .

42CHAPTER 2. INSTRUMENT REQUIREMENTS FROM WEAK LENSING SCIENCES

Figure 2.23: Same as in Fig. 2.21, but the results with near-infrared bands together withthe optical bands. The left panel shows the result with g 0 r0 i0 z 0 + Ks bands, while the rightshows the result with the 7 bands of g 0 , r0 , i0 , z 0 + J, H, Ks . Both panels assume S/N=10 ineach passaband.
The simulation is done in the following manner. First, a mock catalog of galaxies isgenerated by picking ∼ 2200 SED’s randomly from the template SED’s so that the redshift
distribution N (z) of the mock sample reproduces a typical N (z) of galaxies. Here we assume
an N (z) in an analytical form N (z) ∝ z 2 exp(−z/0.37) which roughly represents the N (z) for
IAB < 26 derived from a galaxy photo-z catalog on the Hubble Deep Field North (Furusawa
et al. 2000). This N (z) is consistent with that used in Stabenau et al. (2008). Then,
we add Gaussian noise to magnitudes listed in the mock catalog in each band, i.e., eachgalaxy in the catalog has the same signal to noise ratio (S/N) in each band. Photo-z for two
diﬀerent S/Ns of 10 (typical case) and 30 (very good case) are examined in this simulation.
Finally, we employ a template ﬁtting method involving model spectra derived from Kodama& Arimoto (1997) used in Furusawa et al.(2000) to calculate output photo-z for each galaxyof the input mock catalog.
We show the results of the simulation. First we assume a conventional combination of
passbands g 0 , r0 , i0 , z 0 (SDSS), which is supposed to be employed in the HSC WL survey (see§2.2). Second, we examine addition of one more red band Y at ∼ 1µm to the 4 passbands.
Finally, the case in which we have ideal near-infrared data in addition to the g 0 , r0 , i0 , z 0 bands
is also examined.
The results of the simulation are shown in Figs. 2.21 to 2.23. We also summarize statisticalvalues for the above simulation including photo-z bias and σz of the photo-z in Table 2.2 to
2.4.
From Table 2.2, we see that photo-z determination with g 0 , r0 , i0 , z 0 bands achieve the rms
error (statistical dispersion) of σz = 1.41 at z = 0.0 − 1.5 for all the galaxy sample without
any treatment (i.e., clipping of outliers, see below). However, if we remove catastrophic errors

2.4. PHOTOMETRIC REDSHIFTS

43

Table 2.2: Results of the simulation on accuracies for photo-z for g 0 , r0 , i0 , z 0 passbands. Eachcolumn shows (Col. 1) S/N assumed in each passband, (Col. 2 & 5) bias of mean redshift
calculated by <(photo-z - input-z)/(1+input-z)>; (Col. 3 & 6) σz , rms error of output z
with respect to the input z for the output; and (Col. 4 & 7) fraction of outputs with good
photo-z accuracy of | photo-z - input-z | < 0.5. The values written between parentheses
are the values for samples with a photo-z accuracy of | photo-z - input-z | < 0.5, which arethought to be the values after we eﬃciently remove the catastrophic errors.
z = 0.0 − 1.5bias
σzFraction-0.49 (-0.015) 1.41 (0.13)
48.5%-0.26 (-0.0052) 1.02 (0.090)
70.9

S/N
10
30

z = 1.5 − 3.0
bias
σzFraction
0.045 (0.0010) 1.00 (0.15)
31.7
0.018 (0.0011) 0.72 (0.12)
58.8

Table 2.3: Same as Table 2.2, but for the 5 passbands of g 0 , r0 , i0 , z 0 , Y .
S/N
10
30

z = 0.0 − 1.5bias
σzFraction-0.38 (-0.012) 1.23 (0.11)
65.4-0.14 (-0.0043) 0.73 (0.069)
88.0

z = 1.5 − 3.0
bias
σzFraction-0.011 (0.0015) 0.77 (0.15)
48.2-0.0010 (-0.0015) 0.38 (0.11)
83.4

of | photo-z - input-z | > 0.5, the rms errors can be decreased to σz = 0.13 at this redshift
range. This is the level of acceptable for the WL study (§2.2). This accuracy still seems
slightly worse than present-day sophisticated photo-z work. However, this is because themethod used here does not implement any priors and is not tuned by empirical training-set
approaches. Similarly, the bias of photo-z from all the galaxy sample (<(photo-z - inputz)/(1+input-z)>) is −0.49(= 49%), while after removal of catastrophic samples, it can be
reduced to as small as 1.5%, which is merginally acceptable to the WL studies. The fractionof outputs surviving the outlier-removal is about 50 and 30 percents for 0 < z < 1.5 and1.5 < z < 3 samples, respectively. This seems a great loss, but as was mentioned in §2.2, asigniﬁcant fraction of galaxies can be discarded without severely degrading the cosmologicalparameter accuracy from WL power spectrum, because the weak lensing power spectrum<
3, 000measurement is not shot noise limited as long as only the spectra of multipole range l ∼are included (Jain, Connolly & Takada 2007).

Table 2.4: Same as Table 2.2, but the result in addition of near-infrared bands. S/N is set
to 10 in each band.

band
grizKsgrizJHKsz = 0.0 − 1.5bias
σzFraction-0.11 (-0.011)
0.68 (0.11)
84.2%-0.023 (-0.0056) 0.35 (0.099)
92.5

z = 1.5 − 3.0
bias
σzFraction-0.020 (-0.0044) 0.45 (0.14)
69.5-0.0022 (-0.0020) 0.29 (0.13)
82.0

44CHAPTER 2. INSTRUMENT REQUIREMENTS FROM WEAK LENSING SCIENCES
Thus, with only g 0 , r0 , i0 , z 0 bands, it is possible to archive the photo-z accuracy required
from the HSC WL study if photo-z outliers can be successfully excluded. We are investigatingeﬃcient and robust ways of removing the catastrophic errors. First, it is important to
optimize template SEDs and the ﬁtting method using galaxy samples whose redshifts are
known by spectroscopic observations – the empirical training-set approaches (e.g., Connolly
et al. 1997, Collister & Lahav 2004). We will implement this approach in our photo-z
method. For the training-set, we will use spectroscopic samples from ongoing or planned
surveys such as SDF, SXDS, VIMOS VLT Deep Survey, and BOSS, Second, we will introduce
reasonable priors (e.g., Benı́tez et al. 2000, Stabenau et al. 2008). Importantly, those two
approaches have been already successfully demonstrated by several teams. Third, we may
improve the bias of photo-z by using galaxies with particular spectral types. This uses thefact that because of the distinct 4000 Åbreak, which is an important feature in the SED
in photo-z determination, photo-z for early-type (E) galaxies are more accurately estimatedthan for late-type (Spiral and Disk) galaxies. Similarly, we will introduce a “reliability
parameter” of resulting photo-z utilizing a PDF of photo-z (the likelihood function of thephoto-z obtained from the χ2 -ﬁtting method). A galaxy with a good photo-z tends to havea very peaky PDF, whereas one with a catastrophic photo-z error tends to have a bimodal
or very broad PDF. Such features in the PDF can be used to eﬀectively remove galaxies
with unreliable photo-z.
Although the HSC WL survey is designed to be a stand-alone program with g 0 r0 i0 z 0 band
data, it is interesting to see how additional data improves the photo-z accuracy. This kindof study would be also useful to evaluate a merit of a possible extension of the HSC WL
survey. First, we consider the case that the Y -band data at 1µm is added to the basic
g 0 r0 i0 z 0 data (Table 2.3, Figure 2.22), σz and the bias are improved up to 0.11 and 1.2%
(0.13 and 1.5% for the basic case), respectively. The improvement in the fraction of thegood photo-z galaxies is about 15% (from 50 to 65% for 0 < z < 1.5 and from 32 to 48%
for 1.5 < z < 3). Therefore the improvement in the photo-z accuracy is not very big. We
comment on the expected exposure time for Y -band imaging which is the another factor
to be taken into consideration. Due to the high background, the exposure time of the Y band imaging to a similar depth of other four bands is expected to be about 30 minutes(Y ∼ 25 AB magnitude). Thus adding Y -band results in about 50 percent increase of thetotal exposure time.
Next, we consider the case deep near-infrared data are obtained by a wide ﬁeld NIR
imager on other facilities. The result of the simulation is presented in Figure 2.23 and issummarized in Table 2.4. In this case, we can drastically remove the outliers even without
special treatment such as tuning of template SEDs and employment of priors. Moreover thefraction of accurate photo-z increases up to 92.5% from 48.5% achieved only by the 4 opticalpassbands. Although the deep imaging in NIR bands is very expensive today, even NIR
data in a small area will be useful for a certain purpose described below. Since the photoz from the optical and NIR data is so accurate, it can be used for calibration of photo-z
method. For example, we could supplement the training-set data in the redshift range where
spectroscopic redshifts are poorly obtained (the so-called redshift desert) with the opticaland NIR photo-z. We are investigating this possibility with available optical-NIR data, e.g.,
SXDs and GOODS.

2.4. PHOTOMETRIC REDSHIFTS

2.4.3

45

Estimating the photometric redshift error with clustering

We are investigating a possible method to estimate the photometric redshift error. Thismethod uses the clustering analysis of the photometric redshift sample. Denoting the 3
dimensional mass-power spectrum at the redshift z in real space by Pmass (k, z), the galaxy
power spectrum in the redshift space can be written as¡¢2
P (k, µ, z) = b(z, k) + f (z)µ2 Pmass (k, z)DFoG (k, µ, z)DPZE (k, µ, z),where b(z, k) is the bias factor, DFoG (k, µ, z) is the damping factor due to the Finger of God
eﬀect, µ is the directional cosine between the line of sight direction and the wave number
vector, DPZE (k, µ, z) describes the eﬀect of the photometric redshift error, which can be
written as
¿
·
µ
¶¸À
dχ
dχ 0
DPZE (k, µ, z) =exp iµk∆z − 0 ∆z
dz
dz
¶¸
·
µ
Z Z
dχ 0
dχ
0
0
=
d(∆z)d(∆z )P(∆z)P(∆z ) exp iµk∆z − 0 ∆z,dzdz
where P(∆z) is the distribution function of the redshift error ∆z, and χ = χ[z] is thecomoving distance. In the case of the Gaussian distribution function√£¤P(∆z) = [ 2πσz ]−1/2 exp −∆z 2 /2σz2 ,we have

·
µ ¶ ¸
dχ
2 2
DPZE (k, µ, z) = exp −µ k
σz2 ,
dz

where σz2 is the variance of the error of the photometric redshift. The 3 dimensional (angularaveraged) power spectrum is
√
Z
π
1
1 +1
(3)
2
P (k, z) =dµP (k, µ, z) ' b(k, z) P (k, µ = 0, z).
2 −1
2kσz dχ/dz>
This approximation is valid as long as the variance is large, e.g., σz ∼0.01.
On the other hand, we may deﬁne the 2 dimensional power spectrum, which is essentiallyequivalent to the angular power spectrum, as follows,

P (2) (k, z) ' b(k, z)2 P (k, µ = 0, z).
Thus, the ratio of the 3 dimensional power spectrum to the 2 dimensional spectrum,√P (3) (k, z)
π
1
=
,
(2)
P (k, z)
2 kσz dχ/dz
gives the variance of the photometric redshift error σz . This is a basic idea of the method, butthere will be problems of the non-Gaussian distribution function of the error, and the ﬁnite
range of the redshift distribution of sample, etc. However, the photometric redshift error is
crucial to the precise measurement of the weak lensing statistics. Also it is of fundamental
importance to establish the way of self-calibrating the photo-z error with available data.
The capability of this method must be investigated carefully in future.

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[32] Okabe N., Umetsu K. 2008, PASJ in press[33] Okura Y., Umetsu K., Futamase T. 2007, ApJ, 660, 995[34] Pierre M., Pacaud F., et al. 2006, MNRAS, 372, 591[35] Schneider P. 1998, ApJ, 498, 43[36] Schneider P., Watts P. 2005, A&A, 432, 783[37] Sheth R. K., Tormen G., 1999, MNRAS, 308, 119[38] Simon P. 2005, A&A, 430, 827[39] Simon P., et al. 2007, A&A, 461, 861[40] Sheldon E., et al. 2004, ApJ, 127, 254[41] Stabenau H. F., et al. 2008, arXiv:0712.1594[42] Takada M., Jain B. 2004, MNRAS, 348, 897[43] Takada M., Bridle S. 2007, New. J. Phys., 9, 446[44] Takada M., White M. 2004, ApJ, 601, L1[45] Tasitsiomi A., et al. 2004, ApJ, 614, 533[46] Tegmark M., Peebles P. J. E. 1998, ApJ, 500, L79[47] Umetsu, K., et al. 2004, Mod. Phys. Lett. A, 19, 1027[48] Umetsu K., Takada M., Broadhurst T. 2007, Mod. Phys. Lett. A., 22, 2099
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76, 023504[53] Yoshikawa K., et al. 2001, ApJ, 558, 520

Chapter 3
Hyper Suprime-Cam System
Deﬁnition
In this chapter, we ﬁrstly deﬁne the system of Hyper Suprime-Cam followed by the distribution of the error budget regarding the image quality to each system component. We thendeﬁne the speciﬁcations of those component in terms of image quality.

3.1

System Components

Figure 3.1: Schematic view of planned HSC structure.

51

52

CHAPTER 3. HYPER SUPRIME-CAM SYSTEM DEFINITION

A schematic view of planned HSC system is shown in Figure 3.1. We describe the components of HSC brieﬂy here.
1. Prime Focus Unit (PFU)
Base Frame (BF) The BF is the foundation structure of the entire HSC system andconnects between the telescope inner hub and Stewart platform. The BF will bemade of aluminum and has two slits for the ﬁlter insertion.
Stewart Platform (SP) The SP controls attitude and position of the HSC. The SP
is located between the BF and the top frame. Since each actuator in the SP isbasically a copy of existing one with slight modiﬁcations, it must be reliable.
Top Frame (TF) The TF hangs instrument rotator, dewar frame, dewar, ﬁlter, shutter, lens frame, lens barrel, and lens optics. The TF will be made of aluminumand is located between the SP and instrument rotator.
Instrument Rotator (InR) The InR rotates dewar frame, dewar, shutter, and ﬁlter
in order to compensate the ﬁeld rotation. The InR is located at the joint pointof the TF, the dewar frame, and the lens frame.
Dewar Frame (DF) The DF supports the dewar, electronics, cooler, ﬁlter, and shutter. The DF will be made of aluminum. The stiﬀness of the DF and the lens frame
will be adjusted in order to keep the relative position and attitudes between optics
and detectors.
Lens Frame (LF) The LF connects the lens barrel to ﬁxed side of the InR andsupports the wide ﬁeld corrector. Therefore the lens barrel do not co-rotate. The
LF will be made of CFRP. In order to relieve an internal stress by the diﬀerenceof the coeﬃcient of thermal expansion, invar ring will be installed at the interface
surface to the lens barrel. The LF also has two slits for ﬁlter insertion.
Cable Wrapper (CW) The CW wraps cables and pipings in order to remove slackness in rotation of the InR.
2. Wide Field Corrector (WFC) (see Chapter 4)
Lens Optics Lens corrector at the prime focus of a hyperboloidal mirror.
Lens Barrel (LB) The LB is a structures to hold hold the lens optics.
3. Camera (see Camera / Filter chapter for detail)
Dewar The dewar will hold detectors and keep cool them.
Filter Exchanger The ﬁlter exchanger exchanges ﬁlters.
Filter The ﬁlter will select wavelength pass band of observation.
Shutter The shutter will control exposure on detectors.
There are three major components, PFU, WFC and Camera. We have subcontracted with
Mitsubishi and Canon for the development of the PFU and the WFC, respectively.

3.2. FEASIBILITY STUDY OF THE HSC IMAGE QUALITY

3.2

53

Feasibility study of the HSC image quality

In order to make a feasibility study of HSC image quality, we analyzed the error budget
distribution of the current Suprime-Cam. We were not able to ﬁnd out the exact documentsthat argued the issue comprehensively. Therefore, we collected related material and estimated the values based on the speciﬁcations to each element and others. Table 3.1 shows
the result. Here we deal with the worst case anywhere in the 30 arcmin diameter ﬁeld ofview and the wavelength covering 400 nm to 1000 nm.
FWHM (00 )0.321

Elements
Diﬀraction
OpticsPrimary mirror
Mechanics

FWHM (00 )
0.013
0.180
0.100
0.170

Charge diﬀusion

0.178

Details

FWHM (00 )Optics aber.

0.180

Tracking err.
Setting/Drive err.

0.120
0.120

Table 3.1: Error budget analysis of Suprime-Cam.
If we could realize exactly the same image quality of Suprime-Cam, it would be ideal.
With the goal in mind, we worked on the iterative precess of the feasibility studies of eachcomponent and the error budget re-distribution. Table 3.2 shows the summary of the studies.
The estimated instrument PSF is 0.38300 (FWHM) and this is slightly worse than SuprimeCam case shown in Table 3.1. However, the impact of the deterioration is limited in apractical sense. For example, under the condition of natural seeing of ∼ 0.600 , SuprimeCam will have images of 0.6800 in FWHM of the PSF whereas HSC will have 0.7100 and thediﬀerence is small.

3.2.1

Sensitivity Analysis

In the following, we provide approximate conversion formulae between displacement/tilt of
HSC components and deterioration of image quality are calculated by sensitivity analyses on
the HSC system. We summarize the result of sensitivity analyses in Table 3.3. Coordinate
for the displacements is shown in Figure 3.2. Note that these formulae are calculated using
the lights of entire wavelength range not in a particular band pass. Therefore this is regarded
as the worst case study.

3.2.2

Speciﬁcations of PFU mechanics

We show the speciﬁcations of PFU mechanics component. Mitsubishi has already agreed
upon the speciﬁcation. As described above, the PFU consists of BF, SP, TF, InR, DF, LF,
and CW; accordingly the speciﬁcations of PFU impacts on these elements. We distributed
the budget among following items.
Optics Displacement Optics displacement means collimation errors (shift and tilt) between the WFC and the optical axis and focal point of the primary mirror in all degreeof freedom (DOF), i.e. ∆X, ∆Y , ∆Z, ∆θX , and ∆θY . It goes without saying that∆θZ does not aﬀect on the image quality because each of the WFC and the primary
mirror is practically axisymmetric on Z axis.

54
FWHM (00 )0.383

CHAPTER 3. HYPER SUPRIME-CAM SYSTEM DEFINITION
Elements
Diﬀraction
Optics / ADC

FWHM (00 )
0.013
0.201

Primary mirror
Collimation

0.100
0.169

Tracking

0.155

Detector align.
Dewar ﬂexure

0.084
0.023

Charge diﬀusion
Filter par.
Window par.

0.178
0.050
0.050

Details

FWHM (00 )Optics aber.
ADC drive err.

0.200
0.016

Optics displ.
Rotator axis displ.
Detector tilt
Detector defocus
Stewart platform
Rotator axis runout
Rotator angle err.
Telescope mount
Auto-guider err.

0.063
0.090
0.120
0.046
0.005
0.063
0.068
0.120
0.025

Flexure deviation
Detector displ.
Detector defocus

0.019
0.014
0.009

Table 3.2: Error budget table for the HSC.
Rotator Axis Displacement Rotator axis displacement means the diﬀerence between the
optical axis of the WFC (the center of image rotation) and the mechanical axis of the
InR (the center of instrument rotation). This term can be expressed by using ∆X and∆Y . We present the schematic view of this term in Figure 3.3.
Detector Tilt Detector tilt means the inclination between surface of detectors and the
focal plane formed by the WFC and the primary mirror. This term can be represented
by using ∆θX and ∆θY . We present the schematic view of this term in Figure 3.4.
Detector Defocus Detector defocus means the error on the spacing between the WFC andthe Camera. HSC has no active actuators to adjust the spacing , we should adjust thisspacing adequately at the assembly. The spacing is expressed as ∆Z. We present theschematic view of this term in Figure 3.5.
Tracking Error Tracking error means the image position error in the focal plane due tothe instability of the telescope tracking. This is caused by various reasons including
the smoothness of the SP drive, runout of the rotator axis, rotator drive angle error,
telescope mount error, and auto-guider error.
Now we describe the error budget on each elements / components of the PFU in following
subsections.

3.2. FEASIBILITY STUDY OF THE HSC IMAGE QUALITY
Components
WFC

ADC
CCD

WFC-CCD sep.

55

Displacement Axis Formulae for Deterioration (in 00 ) ID
X
0.466 × ∆X[mm] + 0.0061
1Y
0.463 × ∆Y [mm] + 0.0083
2Z
5.211 × ∆Z[mm] + 0.0148
3
X-rot
14.294 × ∆θX [deg] + 0.0070
4
Y-rot
14.428 × ∆θY [deg] + 0.0047
5Y
0.085 × ∆Y [mm] + 0.0008
6Z
2.703 × ∆Z[mm] + 0.0123
7
X-rot
11.351 × ∆θX [deg] + 0.0116
8
Y-rot
11.352 × ∆θY [deg] + 0.0104
9Z
0.197 × ∆Z[mm] + 0.006
10

Table 3.3: Approximate conversion formulae derived by sensitivity analyses.
3.2.2.1

Optics displacement

We conclude that it is enough to allocate the budget of ∆FWHM = 0.06300 for the optics
displacement. Optics displacement is aﬀected by SP drive error, the repeatability of installation of HSC to the inner hub of the telescope, and telescope ﬂexure repeatability. The
repeatability of installation of HSC depends upon the repeatability of the curvic couplingat the telescope-HSC interface. Therefore we subdivided the budget into these components,∆FWHM = 0.02400 for SP drive error; ∆FWHM = 0.03200 for curvic coupling position repeatability; ∆FWHM = 0.02400 for telescope ﬂexure repeatability.
For the SP drive error, we subdivided the budget 0.02400 into 0.00700 , 0.00900 , 0.01700 , 0.00900 ,and 0.00700 for ∆X, ∆Y , ∆Z, ∆θX , and ∆θY , respectively. We translated these budget tothe speciﬁcations on each axis by using formulae ID 1 to 5 in the Table 3.3 and obtained
values are 2 µm, 2 µm, 0.9 µm, 0.600 , and 0.600 for ∆X, ∆Y , ∆Z, ∆θX , and ∆θY , respectively.
For the curvic coupling position repeatability, we segmented the budget 0.03200 into 0.01100 ,0.01300 , 0.06600 , 0.02100 , and 0.01800 for each axis. We calculate the speciﬁcations on the curvic
coupling in similar way on the case of the SP and results are 10 µm, 10 µm, 10 µm, 3.400 , and3.400 for ∆X, ∆Y , ∆Z, ∆θX , and ∆θY , respectively. In Figure 3.6, we present the schematic
view for the deﬁnition of the repeatability of curvic coupling.
For the telescope ﬂexure repeatability, we use 5 % of the actual result of telescope ﬂexuremeasured on Suprime-Cam as the ambiguity of telescope ﬂexure. We translated these ﬂexures
into the budgets by using formulae ID 1 to 5 in Table 3.3 and results are 0.00800 , 0.01200 ,0.04000 , 0.00700 , and 0.00500 for ∆X, ∆Y , ∆Z, ∆θX , and ∆θY , respectively. The calculated
values do not exceed the distributed budget so we adopt the 5 % values on each direction
as the speciﬁcations. The calculated repeatability of telescope ﬂexures are 3.7 µm, 7.3 µm,
5.7 µm, 0.00400 , and 0.10800 for ∆X, ∆Y , ∆Z, ∆θX , and ∆θY , respectively.

3.2.2.2

Rotator axis displacement∆FWHM = 0.09000 is allocated for the rotator axis displacement. If the optics has no
distortion and auto-guider is working, rotator axis displacement does not aﬀect image quality. Regrettably, the WFC has not distortion-free, the rotator axis displacement should betaken into account to estimate image quality. In the worst case, 180 degree rotation, the56

CHAPTER 3. HYPER SUPRIME-CAM SYSTEM DEFINITION

Figure 3.2: Coordinate system on HSC structure.

displacement of image on the focal plane can be expressed as below:
l ' 8δM ∆r(3.1)

where l is image displacement, δM is distortion parameter at the edge of FOV, and ∆r isrotator axis displacement. Needless to say, l is a function of rotation angle. So the image
length in the case of the rotation angle θR can be represented asl ' 8δM

θR
∆r .
2π

(3.2)

In order to determine the speciﬁcation for the rotator axis displacement, we suppose thatthe rotation angle in one exposure is ∼ 15 degree. Since the image length limit lMAX shouldbe treated as rms image size, ∆FWHM = 0.09000 is converted to the image length limit
of lMAX = ∆FWHM/1.665 × (focal-length) = 4.83 µm. By using of the equation 3.2 and
4.83
× 180
= 207 µm. We divideδM = 0.035 (see WFC chapter), we can get ∆r = 8×0.035
15p22
this ∆r into ∆X and ∆Y satisfying ∆r = (∆X) + (∆Y ) and ∆X ' ∆Y and get the
speciﬁcations for each of ∆X and ∆Y as 146 µm.

3.2. FEASIBILITY STUDY OF THE HSC IMAGE QUALITY

57

Figure 3.3: Schematic view of rotator axis displacement.

3.2.2.3

Detector tilt error∆FWHM = 0.12000 is allocated for the detector tilt. Detector tilt occurs due to the LF tilt,
rotator bearing error, DF tilt, and dewar internal tilt. We subdivided the budget into thesecomponents, ∆FWHM = 0.05500 for LF tilt; ∆FWHM = 0.06400 for rotator bearing error;∆FWHM = 0.06100 for DF tilt; ∆FWHM = 0.06000 for dewar internal tilt.
For the LF tilt, the budget is subdivided into 0.04100 and 0.03600 for ∆θX and ∆θY , respectively. These values are translated into the speciﬁcations by using of the formula ID 8 and
9 in Table 3.3. The results are 9.400 and 8.100 for ∆θX and ∆θY , respectively.
For the rotator bearing error, the budget is subdivided into 0.04600 and 0.04500 for ∆θX and∆θY , respectively. These values are translated into the speciﬁcations by using of the formula
ID 8 and 9 in Table 3.3. The results are 10.800 and 10.800 for ∆θX and ∆θY , respectively.
For the DF tilt, the budget is subdivided into 0.04400 and 0.04300 for ∆θX and ∆θY , respectively. These values are translated into the speciﬁcations by using of the formula ID 8 and
9 in Table 3.3. The results are 10.300 and 10.300 for ∆θX and ∆θY , respectively.
For the dewar internal tilt, the budget is subdivided into 0.04200 for each of ∆θX and ∆θY .
These values are translated into the speciﬁcations by using of the formula ID 8 and 9 inTable 3.3. The results are 1000 and 1000 for ∆θX and ∆θY , respectively.

58

CHAPTER 3. HYPER SUPRIME-CAM SYSTEM DEFINITION

Figure 3.4: Schematic view of detector tilt.

Figure 3.5: Schematic view of detector defocus.

3.2.2.4

Detector defocus error∆FWHM = 0.04600 is allocated for the detector defocus. We will adjust the separation
between optics and detectors by inserting spacers between DF and dewar at the assembly
phase. The thickness of spacers planned to be 200 µm in order to satisfy the error budget on
this component. This examination is based on the conversion formula ID 10 in Table 3.3. For
conﬁrmation, we estimated the image deterioration in a narrow scope of the detector defocus.
The cases of ∆Z < 150 µm are calculated and image deterioration is smaller than 0.0100 for
nominal value of D80 = 0.150√00 in r band (see WFC chapter). In the case of ∆Z = −150 µm,∆FWHM is represented by 0.1582 − 0.1502 /1.523 ' 0.03300 . This value satisﬁed the error
budget 0.04600 . Therefore we concluded that this error budget can be fulﬁlled.
3.2.2.5

Tracking error∆FWHM = 0.15500 is allocated for telescope tracking. As mentioned above, telescope tracking is determined by the smoothness of the SP drive, rotator axis runout in rotation, rotator
rotation angle error, telescope mount, and auto-guider error. We subdivided the budget into
these components, ∆FWHM = 0.00500 for smoothness of the SP drive, ∆FWHM = 0.06300
for rotator axis runout in rotation, ∆FWHM = 0.06800 for rotator rotation angle error,∆FWHM = 0.12000 for telescope mount, ∆FWHM = 0.02500 for auto-guider error.
For the smoothness of the SP drive, the budget is subdivided into X- and Y - direction in
the same weight. Therefore there is 0.00300 ∆FWHM budget for each of them. The image
length should be treated as rms value, so the limit of image length is 0.003/1.665 = 0.00200 .
In order to convert this angle to the length on the focal plane, we multiply this value by the
focal length and obtain 0.2 µm on each direction, ∆X and ∆Y .
For the rotator axis runout, the budget of 0.06300 is subdivided into X- and Y -directionin the same weight. So there is 0.04500 ∆FWHM budget for each of them. Using the same
treatment on the smoothness of the SP drive, we obtain 2.4 µm on each direction, ∆X and3.2. FEASIBILITY STUDY OF THE HSC IMAGE QUALITY

59

Figure 3.6: Schematic view for the deﬁnition of the curvic coupling position repeatability.

∆Y .
An error concerning rotator rotation angle, that is expressed by this equation:l ' Rfp ∆θR(3.3)

where l is image displacement (image length), Rfp is radius of focal plane, and ∆θR is theerror on rotator rotation angle. ∆θR can be estimated easily at 3.100 by using the radius of
FOV of Rfp ' 240 mm
As for the telescope mount, present value of Subaru telescope mounting, ∆FWHM =0.11700 , is used. This value satisﬁes the budget 0.12000 .
The auto-guider position error can be directly converted to the image quality and thespeciﬁcation on this term is 1.5 µm on each direction. This means we should detect theposition of guiding star at the accuracy of 0.1 pixel in X- and Y - direction on the guiding
CCD.

3.2.3

Speciﬁcations for WFC and ADC

We allocate 0.20000 for WFC optics (see Chapter 4). The atmospheric dispersion compensator(ADC) in the WFC is moved in order to compensate the atmospheric dispersion because
the dispersion depends upon the elevation of the telescope. ADC drive mechanics will be
constructed to meet the speciﬁcation, ∆FWHM = 0.01600 . In this case, ADC drive error canbe estimated at 173.6µm by using of the formula ID 6 in Table 3.3.

3.2.4

Speciﬁcations for ﬁlter

We budget 0.05000 for ﬁlter. Because thickness variability of the ﬁlter makes a tilt of focal
plane, parallelism of ﬁlters should be less than 4000 .

3.2.5

Speciﬁcations for Camera mechanics and sensors

3.2.5.1

Surface ﬂatness of detectors

Error budget for detector surface ﬂatness is ∆FWHM = 0.08400 . Detector surface ﬂatness isaﬀected by cold plate (CP) surface ﬂatness and detector thickness uniformity. The surface
ﬂatness of the CP consists of the manufacturing tolerance and the warp by external force,so we subdivided the budget into these components, ∆FWHM = 0.04100 for CP surface
manufacturing tolerance; ∆FWHM = 0.02700 for CP surface warp; ∆FWHM = 0.06800 for
detector thickness uniformity.
For these components, the conversion formulae is the ID 7 in Table 3.3 is used to translate
into the speciﬁcations. The budgets 0.04100 , 0.02700 , and 0.06800 for CP surface manufacturing

60

CHAPTER 3. HYPER SUPRIME-CAM SYSTEM DEFINITION

tolerance, CP surface warp, and detector thickness uniformity are converted to 15 µm, 10 µm,
and 25 µm, respectively.
3.2.5.2

Charge diﬀusion on detectors

Since the FDCCD has thick Si layer in it, charge diﬀusion eﬀect in CCDs can not be disregarded. We examined the charge diﬀusion eﬀect on a pilot type CCD and F/2 beam, andfound a parameter of charge diﬀusion σD ' 7 µm at 200 µm thickness. This radius can beconverted to the ∆FWHM on the focal plane and the result is ∆FWHM = 0.17800
3.2.5.3

Flexure ﬂuctuation in the dewar

Since the rigidity of cold surface support is not symmetrical, ﬂexure/displacement of cold
surface caused by gravity force will change with the rotator angles. ∆FWHM = 0.01900 isallocated for this ﬂexure ﬂuctuation. We divide the ∆FWHM into X- and Y - direction andobtain 0.01300 for each direction. As mentioned section 3.2.2.5, the image length should betreated as rms value and the limit is 0.013/1.665 = 0.00800 . This angle can be converted to
the length on the focal plane by multiplying the focal length and we obtain the value 0.7 µm
for the speciﬁcations of the ﬂexure ﬂuctuation on X- and Y - axis.
3.2.5.4

Flexure in the dewar∆FWHM = 0.00900 is allocated for the X- and Y - axes displacement of the cold plate. The
X- and Y - axes ﬂexure in the dewar can be treated as the rotator axis displacement. The
budget 0.00900 can be converted to the speciﬁcations on ∆X and ∆Y in the same way of thecase of rotator axis displacement, and result is 15 µm on each directions.
∆FWHM = 0.01400 is allocates for the Z- axis displacement of the cold plate. The Zdisplacement can be treated in the same way of detector defocus error term. The budget
0.01400 can be converted to the speciﬁcations on ∆Z in the same way of the case of detector
defocus error, and result is 40 µm.
3.2.5.5

Parallelism of dewar window

We budget 0.05000 for the parallelism of dewar window. The identical situation to those on
ﬁlter, parallelism of dewar window should be less than 4000 .

3.2.6

Summary

We summarize the speciﬁcations for components in Table 3.4.

Auto-guider

TelescopeSensors

ADC drive
WFC
Filter
Dewar

Frame

RotatorSP

Components

drive error
smoothness
axis displacement
axis runout
angle error
total tilt
(LF)
(Rotator bearing)(DF)
WFC-CCD sep.
Curvic coupling
drive error
optical aberration
parallelism
CP surface ﬂatness
CP warp
CP tilt
CP displacement (X,Y)
CP displacement (Z)ﬂexure deviation
window parallelism
thickness unif.
charge diﬀusion
ﬂexure repeatability
mounting error
position error
1.5

5.7

25

40

15
10

0.004

10

3.4

17.6
9.4
10.8
10.3

∆θX
[00 ]
0.6

0.108

10

3.4

17.0
8.1
10.8
10.3

∆θY
[00 ]
0.6

3.1

∆θR
[00 ]

173.6

∆ADC
[µm]

Table 3.4: Summary of the speciﬁcations.

1.5

7.3

1

1

3.7

15

10

200
10

∆Y
∆Z
[µm] [µm]
2.0
0.9
0.2
146
2.4

15

10

∆X
[µm]
2.0
0.2
146
2.4

40

40

k
[00 ]

7

0.046
0.032
0.016
0.200
0.050
0.041
0.027
0.060
0.009
0.014
0.019
0.050
0.068
0.178
0.044
0.120
0.023

σD ∆FWHM
[µm]
[00 ]
0.024
0.005
0.090
0.063
0.068
0.104

3.2. FEASIBILITY STUDY OF THE HSC IMAGE QUALITY
61

Chapter 4
Wide Field Corrector (WFC)
We started feasibility studies of the wide ﬁeld corrector (FOV > φ 1.5◦ in 2002 in collaboration
with the late Dr. Takeshi, who designed the current Suprime-Cam corrector at Canon [3, 4].
This was realized by the understandings and supports of a optics ﬁrm, Okamoto Optics in
Yokohama. Based on the studies Komiyama et al. (2004)[5] discussed a possible opticaldesign for φ 2 deg option and they further improved the design later[6] to have more feasible
manufacturability by reducing the asphericity. They also discuss the φ 1.5 deg design, andthese satisfy our design goal, and are plausible.
There is an independent line of optical design study that originally came from the WFMOS
feasibility study. Peter Gillingham, who used to work on the feasibility study at AngloAustralian Observatory, and now joins the HSC team, has improved the original design
signiﬁcantly to satisfy both HSC and WFMOS optical requirements. HSC requires smaller
images but does not require as wide wavelength coverage as WFMOS; WFMOS allows larger
ﬁeld curvature than HSC does and so on. One of the plausible design was shown in a SPIE
paper[7].
In 2006, after the budget of HSC development was partially funded, we started to workwith Canon formally. This chapter summarizes the results of the design studies made incollaboration with Canon.

4.1

WFC Requirements

The requirement for each ﬁlter is described in Table 4.1. Although the system deﬁnitionis given by FWHM (Chapter 3), the requirement is speciﬁed by 80 % encircled Energy
diameter, D80 , at each ﬁlter here because of some historical reasons. If the gaussian proﬁleis assumed, the conversion is the following,
D80 = 1.52 FWHM

(4.1)

For the evaluation of the optics performance, three wavelength are chosen to represent
each ﬁlter which are shown in Table 4.1. These three wavelengths have equal weight in thecalculation. Also these requirements must be met from the telescope Elevation (EL) = 90
degree all the way down to EL = 30 degree. Generally, the image quality tends to deteriorate
more when the elevation becomes lower. Observing log book of Suprime-Cam for nearly 10
years tells that 95 % of observation is done between EL of 30 and 85 degree. For this reason
we only care about the performance at EL > 30 degree.
63

64

CHAPTER 4. WIDE FIELD CORRECTOR (WFC)

g ﬁlterr ﬁlter
i ﬁlter
z ﬁlter
y ﬁlter< 0.5”< 0.3”< 0.3”< 0.3”< 0.4”(420, 470, 530nm)(570, 620, 670nm)(710, 760, 820nm)(870, 910, 960nm)(970, 1020, 1070nm)

Table 4.1: D80 (Encircled energy) requirements

These ﬁgures, FWHMWFC are the total budget including design itself, manufacturing andassembly errors.
The error budget for designed value, FWHMdesign is shown in Table 4.2.

r,i,z ﬁlter
g,y ﬁlter

FWHMdesign< 0.132”< 0.164”

D80design< 0.2”< 0.25”

Table 4.2: Error budget for HSC WFC design
The error budget for fabrication, assembly and environmental change, FWHMfab is therefore set at < 0.145”(r,i,z ﬁlter) to meet the total budget constraint where
FWHM2WFC = FWHM2design + FWHM2fab(4.2)

In addition to these requirements WFC must meet with restrictions from mechanical system,
such as lens diameter sizes,lens system total length ,image plane position and total weight.

4.2

WFC Design

Here we show the latest design that meets the speciﬁcations.

4.2.1

Optics Layout

WFC optical system consists of 7 lenses (7 elements)[1], of which 2 lenses (2 elements)
work as the Atmospheric Dispersion Corrector (hereinafter called ADC). Those lenses are
made of Fused Silica, BSL7Y and PBL1Y (See Figure 4.1). As described in the following
section, these glasses which sizes are over 600mm diameter can be manufactured by someglass companies.
All lenses have an aspherical surface, which are so eﬀective that aberrations are compensated. In fact, some of the surface adopt the convex-aspheric which is usually avoided becauseof the complexity of the measurement setup. Because Canon has an original metrology andapparatus to characterize the lens surface, the design can incorporate such freedom. This is
one of the features of this WFC design. The ﬁrst lens’s eﬀective diameter is 820 mm andtotal lens system length is 1845.3 mm. This WFC optical system qualiﬁes the mechanicalrestrictions and weight restriction.

4.2. WFC DESIGN

65

The plane parallel glasses behind the ﬁfth lens are the ﬁlter and the dewar window. Thisﬁlter consists of 3 plane parallel of 7 mm thickness each (see Chapter 7). The thickness ofthe dewar window is 37 mm. These are all Fused Silica.
General optical data is shown in Table 4.3. This image size is larger than paraxial image
size (about 480 mm) because of the distortion (+3.2 %).
focal length
18320mmimage scale
0.888[mm/arcsec]
image size(1.5deg) 495mm(diameter)
Table 4.3: General optical data

4.2.2

Designed Performance

In this section, we verify the designed performance of the optics.
4.2.2.1

Encircled Energy Diameter

Figure 4.2 ∼ Figure 4.5 show the growth curve of the encircled energy as a function of radius.
i and g ﬁlter at EL of 90 and 30 degree are calculated. We examine the ﬁeld positions listed
in Table 4.4. The maximum value of horizontal axis is set at 0.3/2” which is the total error
X-Field(degree)
0
0
0
0
0
0
0
0
0
0
0

Y-Field(degree)
00.1875
0.375
0.525
0.675
0.75
-0.75
-0.675-0.525-0.375-0.1875

Table 4.4: the ﬁeld positions
budget including manufacturing error.
D80 of each ﬁlter[1] is listed in Table 4.5. The performance qualiﬁes the requirements inall the ﬁlter pass band.
4.2.2.2

Comparison with Suprime-Cam’s WFC

We compare the designed performance with Suprime-Cam’s one. Table 4.6 and Table 4.7
show the comparison of D80 in each ﬁlter at the elevation of 90 and 30 degree, respectively.

Figure 4.1: WFC Layout

G1

*

*

G2

ADC1

ADC2

1845.3mm

G3

*

G4

*

G5

*

* : Aspherical Surface

400mm

Filter+ Dewar window

191.6mm

Glass list
G1
SILICA
G2
BSL7Y
ADC1
BSL7Y
ADC2
PBL1Y
G3
PBL1Y
G4
BSL7Y
G5
SILICA

66
CHAPTER 4. WIDE FIELD CORRECTOR (WFC)

410mm

4.2. WFC DESIGN

67

Figure 4.2: i-ﬁlter EL = 90

Figure 4.3: i-ﬁlter EL = 30

68

CHAPTER 4. WIDE FIELD CORRECTOR (WFC)

Figure 4.4: g-ﬁlter EL = 90

Figure 4.5: g-ﬁlter EL = 30

4.2. WFC DESIGN

69

g-ﬁlter
r-ﬁlter
EL = 90 i-ﬁlter
z-ﬁlter
y-ﬁlter

ADC-shift(mm) D8080 (”)
0
0.163
0
0.150
0
0.146
0
0.151
0
0.199

g-ﬁlter
r-ﬁlter
EL = 30 i-ﬁlter
z-ﬁlter
y-ﬁlter

19.8
21.0
19.8
10.4
4.2

0.196
0.187
0.191
0.186
0.216

requirement
0.25”
0.2”
0.2”
0.2”
0.25”
0.25”
0.2”
0.2”
0.2”
0.25”

Table 4.5: D80 at each ﬁlter

Two tables clearly show that HSC WFC’s performance is almost equivalent with that ofSuprime-Cam’s WFC.

0.00
g ﬁlter 0.11
0.10
r ﬁlter 0.09
0.10
i ﬁlter 0.11
0.08
z ﬁlter 0.15
0.06

0.125 0.25 0.5 0.75 [deg]0.13 0.18
0.10 0.11 0.15 0.16
0.07 0.09
0.11 0.12 0.14 0.15
0.11 0.11
0.09 0.12 0.14 0.13
0.15 0.15
0.07 0.11 0.15 0.15
upper:Suprime-Cam, lower:HSC

Table 4.6: D80 at EL = 90 in arcsec

Since the focal length is 15 % longer in HSC case, the plate scale has changed. Table 4.8
and Table 4.9 represent the result in the unit of 15 µm pixel for direct comparison.

4.2.2.3

U-band performance

Although the original WFC requirement does not include u-band performance, some observers might be interested in the expected performance. Figure 4.6 and Figure 4.7 show
the results examined at the wavelength of 350, 370 and 390 nm with equal weight. This
wavelength coverage is only a part of formal U-band where we expect suﬃcient transmission
eﬃciency on our WFC (See Figures 4.28 and 4.29). The designed D80 is more than twice aslarge as those of longer wavelength.

70

CHAPTER 4. WIDE FIELD CORRECTOR (WFC)-0.75 -0.525
g ﬁlter
0.20
r ﬁlter
0.18
i ﬁlter
0.17
z ﬁlter
0.19-0.25 -0.125 0.00 0.125
0.21
0.15 0.14 0.15
0.14
0.19
0.16 0.14 0.14
0.11
0.08 0.07 0.12
0.15
0.18
0.15 0.14 0.14
0.13
0.12 0.12 0.13
0.16
0.17
0.13 0.13 0.13
0.16
0.16 0.16 0.15
0.16
0.15
0.11 0.11 0.11
upper:Suprime-Cam, lower:HSC

0.25
0.19
0.12
0.13
0.12
0.11
0.12
0.12
0.12

0.525

0.75

0.20

0.19

0.17

0.15

0.17

0.14

0.17

0.16

[deg]

Table 4.7: D80 at EL = 30 in arcsec

0.00
g ﬁlter 0.53
0.60
r ﬁlter 0.43
0.59
i ﬁlter 0.55
0.45
z ﬁlter 0.73
0.36

0.125 0.25 0.5 0.75 [deg]0.64 0.91
0.61 0.67 0.87 0.97
0.36 0.44
0.62 0.71 0.82 0.88
0.53 0.56
0.53 0.68 0.83 0.76
0.72 0.73
0.42 0.62 0.86 0.89
upper:Suprime-Cam, lower:HSC

Table 4.8: D80 at EL = 90 in a 15 µm pixel-0.75 -0.525
g ﬁlter
1.16
r ﬁlter
1.05
i ﬁlter
1.03
z ﬁlter
1.09-0.25 -0.125 0.00 0.125
1.03
0.76 0.68 0.72
0.80
1.13
0.93 0.85 0.85
0.52
0.37 0.35 0.59
0.89
1.08
0.88 0.83 0.81
0.63
0.60 0.61 0.63
0.92
1.01
0.79 0.77 0.76
0.80
0.77 0.80 0.73
0.96
0.88
0.63 0.63 0.65
upper:Suprime-Cam, lower:HSC

0.25
0.93
0.73
0.63
0.71
0.56
0.71
0.57
0.73

0.525

0.75

1.16

1.11

1.02

0.91

1.01

0.84

1.01

0.97

Table 4.9: D80 at EL = 30 in a 15 µm pixel[deg]4.2. WFC DESIGN

71

0.5″
D80 = 0.428″
Figure 4.6: u-ﬁlter EL = 90

0.5″
D80 = 0.455″
Figure 4.7: u-ﬁlter EL = 30

72
4.2.2.4

CHAPTER 4. WIDE FIELD CORRECTOR (WFC)
Vignetting and Distortion

Because the size of the ﬁrst lens is limited we have vignetting in most of the ﬁeld of view.
Figure 4.8 shows the vignetting. Maximum vignetting occurred at the edge of the ﬁeld is
25.6 %.

Figure 4.8: Vignetting

Distortion plot at i ﬁlter is shown in Figure 4.9 and this is calculated only at 760nm (middle
wavelength of i ﬁlter). Distortion depends slightly on the wavelength and Table 4.10 shows
the maximum value in each ﬁlter. These are about 3.2 % which is about 3 times larger thanthat of Suprime-Cam (0.9 %).
Fitting the distortion with the 4th order polynomial yields,
f (y) = 1.77y 4 + 4.63y 2 ,where y is the ﬁeld angle (See Figure 4.10).

g ﬁlter 3.181%
r ﬁlter 3.186%
i ﬁlter 3.184%
z ﬁlter 3.181%
y ﬁlter 3.177%
Table 4.10: Maximum distortion at each ﬁlter(4.3)

4.2. WFC DESIGN

73

0.8
0.7
0.6

angle

0.5
0.4
0.3
0.2
0.1
0
0

0.5

1

1.5

2

2.5

3

3.5

3

3.5

distortion

Figure 4.9: distortion plot at 760nm(i ﬁlter)

0.8
0.7
0.6

angle

0.5
0.4
0.3
0.2
0.1
0
0

0.5

1

1.5

2

2.5

distortion
Desined valuea * y^2 + b * y^4(a = 4.63, b = 1.77)

Figure 4.10: ﬁtted distortion

74
4.2.2.5

CHAPTER 4. WIDE FIELD CORRECTOR (WFC)
Performance beyond φ = 1.5 degree

We plan to use the peripheral ﬁeld for auto guiding and focusing. Some sensors are locatedoutside φ = 1.5 degree ﬁeld of view. In order to know the deterioration of the performance
beyond 1.5 deg, we examine the optics performance at 1.6 degree in this section. Table 4.11
and Figure 4.11 show the D80 and vignetting diagram, respectively. Beyond 1.5 degree thevignetting suddenly increases and reaches 39.5 % at 1.6 degree.
EL = 90 degree EL = 30 degree
u ﬁlter
0.560
0.621
g ﬁlter
0.315
0.382
r ﬁlter
0.234
0.276
i ﬁlter
0.191
0.226
z ﬁlter
0.179
0.212
y ﬁlter
0.242
0.262
Table 4.11: D80 at 1.6 degree

0.8°
Figure 4.11: Vignetting up to 1.6 degree

4.3. ATMOSPHERIC DISPERSION CORRECTOR

4.3

75

Atmospheric Dispersion Corrector

Atmospheric dispersion makes an image performance worse at the low elevation. Additional
optics is required to correct the dispersion, ADC. A conventional type of ADC is invented
by Epps[2]. This type has 2 prisms and rotate these prism in opposite direction to eachother to compensate the eﬀect of atmospheric dispersion. But this ADC requires 4 optics
elements in total.
Since the weight constraint is so tight that we adopted the lateral shift type ADC[3],[4]invented by the late Dr.Takeshi for the Suprime-Cam WFC. This ADC only requires twoelements and at the same time, these work as doublet (achromatizing lens) in optical system.
This enabled the optics system as compact as possible. The brief description how this ADC
work is shown in Figure 4.12.

Figure 4.12: The lateral shift ADC consists of two glasses whose index are the same forthe primary wavelength. The ﬁrst and the third surface are made ﬂat and the second has
a curvature. Therefor, this ADC works as plane parallel for the primary wavelength but
for other wavelengths it works as doublet (left panel). As an example, the one ray of light(not primary wavelength) is refracted downward. Actually this refraction direction depends
on wavelength and glass index of these lenses. Then shift this ADC along the direction
perpendicular to optical axis as shown in the right panel. Non primary light is refracted
upward because the incident angle (the angle between the ray angle and the normal of thesecond surface) is changed by shift while the direction of the primary wavelength remains
the same.
The actual eﬀect is as the following. Top panel of Figure 4.13 is the growth curve of theencircled energy and top panel of Figure 4.14 is spot diagram if ADC is not activated at
EL=30 degree. These ﬁgures show that image performance is deteriorated by atmosphericdispersion quite seriously. Bottom panels of Figure 4.13 and Figure 4.14 shows improvement
thanks to the ADC function.
Because of the limited availability of the large glass blank, the selection of glasses for
the ADC is not ideal. This leads to the elevation dependence of the distortion parameters.

76

CHAPTER 4. WIDE FIELD CORRECTOR (WFC)

Figure 4.13: Encircled energy growth curve when ADC is NOT in use (top) at the elevationof 30 degree and with the ADC activated (bottom). r-band case is considered here.

4.3. ATMOSPHERIC DISPERSION CORRECTOR

77

Figure 4.14: Spot diagram obtained when ADC is NOT in use at the elevation of 30 degree(top) and with the ADC activated (bottom). r-band case is examined here.

78

CHAPTER 4. WIDE FIELD CORRECTOR (WFC)

There is the index diﬀerence between BSL7Y and PBL1Y which are glasses for ADC,then thisdiﬀerence makes a ray of light refracted more widely compared to Suprime-Cam. Actually
the index diﬀerence of HSC ADC is 0.03274. On the other hand, Suprime-Cam ADC glasses
are S-BSM14 and S-TIM5, so the index diﬀerence is 0.0017, which is very small. Therefore if
HSC ADC is shifted and the refracted angle is changed, the ray of light is focused at deﬂected
position, which looks like a distortion. The angle diﬀerence between the chief rays of EL =90◦ and EL = 30◦ at the last surafec of ADC is about 0.057◦ for HSC while this value isabout 0.0023◦ for Suprime-Cam. This angle diﬀerence for HSC is over 20 times as bigger as
Suprime-Cam’s value. This eﬀect has to be taken care of during the data processing.
Two elements of the ADC are not cemented because it is diﬃcult to make it due to thelarge size (∼ 600 mm). Therefor the ADC has an air space in between which could cause aghost. As a result of ghost analysis shown below, it is not so serious.

4.4

Ghosting analysis

We have carried out the calculation of ghosting[1],[8] with all the possible combination
between two surfaces. We neglect reﬂections over 4 times because these multiple reﬂection
rays hardly to come into image surface and the intensity is very small because of the small
reﬂection eﬃciency of the lens.
It was very hard to set the design constraint regarding the level of ghosting. Therefore, we
use the Suprime-Cam design as a reference and examine whether HSC ghosting is equivalent
with that that of Suprime-Cam.
The most visible ghost of Suprime-Cam is shown in Table 4.12 where the level of ghost
is characterized by the following values; DISC is a ghost spot radius at image plane forcentral object (angle 0). When DISC value is small, ghost rays strongly concentrates on
the image plane. Therefore bigger DISC value is preferred. Illuminance ratio is the ratio ofa ghost and a star image illuminance. This value is calculated as following way. We supposethat the star image size is 0.02 mm and the luminous intensity of this image is 1, so the star
illuminance is 1/0.022 π. Also we suppose that the reﬂectance of lens surface is 0.01 and thereﬂectance of ﬁlter and CCD surface is 0.1. If the reﬂection between one lens surface and one
ﬁlter surface, the luminous intensity is 1 × 0.01 × 0.1. And if this ghost DISC is described as
DISCghost , the illuminance of this ghost image is described as 1 × 0.01 × 0.1/π(DISCghost )2
Then we get the equation between illuminance and DISC (See Equation 4.4).

DISC
5.59mmIlluminance Ratio 1.1×10−7
Table 4.12: The parameters of the most visible ghost of Suprime-Cam’s WFC( the reﬂection
between ﬁlter surfaces and set the star size 0.5arcsec (φ0.0371mm))

Equation 4.4 shows the relation between these values. this value shows directly the ghost
intensity and the small value is preferred.
0.022
× η1 × η2 ,Illuminance Ratio =DISC 2(4.4)

4.4. GHOSTING ANALYSIS

79

where η1 and η2 is the reﬂectance of the ﬁrst and the second surface, respectively.
When we consider reﬂections between two lens surfaces, the result is as the following Table
4.13.

reﬂection surface1 reﬂection surface2 DISCADC2R1ADC1R2
3.234
G3R1
G2R2
6.175

reﬂectance1 reﬂectance2
0.01
0.01
0.01
0.01

illuminance ratio
3.824E-09
1.049E-09

Table 4.13: the list of ghosts between lens surfaces

The most visible ghost occurs at the reﬂection between ADC2R1 and ADC1R2. ADC2R1
is the ﬁrst surface of the convex lens of ADC and ADC1R2 is the second surface of theconcave lens of ADC. Because illuminance ratio is smaller the impact is smaller than theSuprime-Cam’s worst case.
Hereinafter we use some abbreviation,such as ADC1R2 and G2R1. The anterior part ofthis abbreviation,such as ADC1 and G2,means which element is. See the Figure 4.1. And
F1∼F3 means ﬁlters,then F1 means the ﬁrst ﬁlter and F2 means the second ﬁlter. And theposterior part means which surface is, then R1 means the ﬁrst surface of a element and R2means the second surface of a element.
When we consider the reﬂection between a lens surface and a ﬁlter surface, the worst caseoccurs at the reﬂection between F2R2 and G4R1 as is shown in Figure 4.15. The result issummarized in Table 4.14.

reﬂection surface1 reﬂection surface2 DISC
F2R2
G4R1
1.288
F3R1
G4R1
2.675
F2R1
G4R1
4.467
F3R2
G4R1
5.855
F1R2
G4R1
8.430

reﬂectance1 reﬂectance2
0.1
0.01
0.1
0.01
0.1
0.01
0.1
0.01
0.1
0.01

illuminance ratio
2.413E-07
5.588E-08
2.004E-08
1.167E-08
5.628E-09

Table 4.14: the list of ghosts between a ﬁlter surface and a lens surface

This ghost is more visible than Suprime-Cam’s. When we made the calculation, the threeglass of the ﬁlter is equally spaced. We noted that by tuning the position of each glass theghost size can be changed. As a results of studies we found a better location to minimize theilluminance ratio down to 5.4 × 10−8 , which is smaller than Suprime-Cam’s the most visible
case.
Note that the ghost can be vanished by the median stacking of dithered images because
the position of ghost changes as telescope is slewed. So impact is limited. The position of theghost occurred between ﬁlter surface and CCD, however, nearly coincides with the original
image on the focal plane. This cannot be eliminated by the median stacking procedure. In
this case, we usually have to mask the neighbors of bright stars which results in the reductionof a ﬁeld of view. We further study the ghost of this case in Chapter 7.

80

CHAPTER 4. WIDE FIELD CORRECTOR (WFC)

Figure 4.15: Reﬂection between F2R2-G4R1

4.5. AVAILABILITY OF GLASS BLANK

4.5

81

Availability of glass blank

Table 4.15 presents requirements for the glass blank[9].

Striaeno visible striae
Birefringence≤ 5nm/cm
Bubble and foreign material total cross section of bubbles(mm2 /100 ml) ≤ 0.1 ∼ 0.25
Refractive index±0.00050 uniformity
Abbe number±0.5% uniformity
Table 4.15: required qualities

The glass speciﬁc requirements and the availability are shown below.
Fused silica
Eﬀective diameter is over 820 mm and the homogeneity ≤ 5 ppm except for G1 that requires≤ 3 ppm. This glass blank are available from Shin-etsu Quartz or Corning.
BSL7Y
Eﬀective diameter is over 630 mm and the homogeneity ≤ 5 ppm. This glass blank areavailable from Ohara or Schott.
PBL1Y
Eﬀective diameter is over 610 mm and the homogeneity ≤ 5 ppm. This glass blank are alsoavailable from Ohara or Schott.

4.6

Design of Lens Barrel Assembly

Each lens element is retained by each lens frame which is made of cordierite which features
a low coeﬃcient of thermal expansion and a high Young’s modulus as is summarized inTable 4.16.
The lens frames are stacked and formed the lens barrel assembly. The schematic view of
the assembly[10] is shown in Figure 4.16.

Coeﬃcient of Thermal Expansion
Young’s modulus
Density< 0.1 ppm
∼ 140 Gpa
∼ 2.7 kg/m3

Table 4.16: Physical property of cordierite

82

CHAPTER 4. WIDE FIELD CORRECTOR (WFC)

Total weight 893.0kg

Primary Mirror
Focus

Image plane

Last lensFigure 4.16: Layout of the lens barrel assembly4.7. IMPACT OF LENS BARREL FLEXURE ON THE IMAGE QUALITY

83

In order to qualify the tight weight constraint, the cordierite is the only material that
we can adopt. However, the drawback is its brittleness. It calls for maximum cautions forhandling. We discuss the additional countermeasures to secure the safety handling operationin Chapter 8.

4.7

Impact of Lens Barrel Flexure on the Image Quality

We simulate the deterioration caused by the deformation due to its own weight. This simulation includes the deﬂections of lens positions and deformations of lens surfaces. The Lens
Frame that holds the barrel is also included in the simulation and the enforced displacement,
estimated from an independent FEM analysis, is applied on the inter-facial boundary withthe rest of the camera system. Figure 4.17 shows exaggerated view of the deformation[12].

Figure 4.17: Lens barrel deformation

The middle panel in Figure 4.18 shows the degree of deterioration of the image quality
which turns out to be ∼ 20 %. We noticed that the deterioration can be recovered if we adjust
the position and attitude of the lens and the camera system using the hexapod actuators.
The bottom panel in Figure 4.18 shows the result of recovery which is indeed equivalent withthe original image quality. We will implement the attitude adjustment parameter into our
new control system of the prime focus unit.

84

CHAPTER 4. WIDE FIELD CORRECTOR (WFC)

NominalUnder deformation

Adjust focus & lensposition by Hexapod
Figure 4.18: Encircled energy growth curve of the deformed lens and the assembly due toits own weight (middle). Elevation of 30 degree and r band ﬁlter is considered here. After
adjusting the attitude of the HSC, the most of the deterioration is recovered (bottom panel).

4.8. IMPACT OF THE TILT AND THE DISPLACEMENT OF CCD WITH RESPECT TO THE LENS

4.8

Impact of the tilt and the displacement of CCD
with respect to the lensWe evaluate the deterioration of image quality caused by the CCD tilt and the displacementwith respect to the WFC for r ﬁlter at EL = 30 because this is the most severe case. Because
the ADC is activated at EL = 30, the optical system is not symmetric about the optical axis.
Therefore, we consider both +20 sec and −20 sec CCD tilt as is shown in Figure 4.19. The+20 sec was the provisional value in the error budget consideration. Here, we try to estimate
the impact of mechanical misalignment of WFC and CCD because we do not implement any
active mechanism to adjust the relative position between the lens barrel and CCD dewar.

WFC

CCD

WFC

CCD

WFC

CCD

Figure 4.19: Deﬁntion of CCD tilt and shift

CCD tilt
Left panels in Figures 4.20 and 4.21 shows deterioration caused by +20 sec and −20 sec tilt,
respectively[13]. The full scale of the horizontal axis is equivalent with D80 of 0.3”. Rightpanels in these ﬁgures shows the degree of recovery after the adjustment of attitude of thecorrector-CCD system using the hexapod. It is interesting to note that deterioration on the+ side is mostly recoverd the original performance after the correction where as on the −side, the recovery is almost halfway.

86

CHAPTER 4. WIDE FIELD CORRECTOR (WFC)

0.3″
Figure 4.20: CCD tilt 20sec (left) CCD tilt 20sec with adjust (right)

Figure 4.21: CCD tilt -20sec (left) CCD tilt -20sec with adjust (right)

4.9. TOLERANCE ANALYSIS87

CCD displacement along the optical axis

14

0.156

13.8

0.154

13.6

0.152

13.4

0.15

13.2

0.148

13

0.146

12.8

0.144

12.6

D80(um)

14.2

0.158

D80"
D80(um)

10
0
15
0
20
0

14.4

0.16

0

14.6

0.162

50

0.164

-2
00
-1
50
-1
00
-5
0

D80″

The image performance deterioration caused by CCD displacement could be compensatedby using the hexapod. Figure 4.22 shows the recovered D80 (r-ﬁlter at EL = 90) for a given
displacement from -200 to +200 µm. It turns out the mechanical tolerance along the opticalaxis is quite large and ∼ 200 µm is allowed.

Δd(um)

Figure 4.22: Recoverd D80 for a give CCD displacement. Left-side y axis is shown in arcsecond whereas right-side y axis is in µm.

4.9Tolerance analysis

We carried out the tolerance analysis by taking fabrication, assembly and other errors into
account[14]. Table 4.17 is summarizes the sources of the error and their values.
The ﬂow chart of the analysis is shown below:(1) Giving the worst value of glass index error⇓(2) Optmization of curvatures, aspherical coeﬃcients,lens thicknesses
and lens spaces (Recovery to the original design value is expected)⇓(3) In addition to (2) lens data, Fabrication errors (curvatures, lens
thicknesses, surface tilt,glass homogeneity and surface accuracy errors)
added on by Monte Carlo method. Then choose the worst lens data.
⇓(4) Optimization spacing of lens⇓(5) In addition to (4) lens data, Assembling errors (lens tilt,lens decenterand residual spacer errors) added on by Mote Carlo method.
Then choose the worst lens data and adjust focus and ADC shift.
⇓

88

CHAPTER 4. WIDE FIELD CORRECTOR (WFC)

Fabrication error
Surface accuracy
Curvature
Glass index homogeneity
Glass index error
Lens thickness
Surface tilt

25 nm RMS
±10 fringes (λ = 633 nm)
5 ppm (2 ppm for G1 and G5)±0.00050± 0.2 mm
± 15 sec(±10sec for G1 and G2)

Assembling error
Decenter
Tilt
Space (residual spacer error)± 0.05mm (± 0.03mm for G1 and G2)± 15 sec (±10 sec for G1 and G2)± 0.01 mm

Environmental change
Surface deform under gravity
result of FEM analysis at EL = 30
Change of space,decenter and tilt under gravity result of FEM analysis ad EL = 30
Table 4.17: Considered errors in the tolerance analysis(6) Enforce the deﬂection of lens shift and tilt and deformation of lens surface(Figure 4.24)⇓(7) Adjust focus and ADC shift(Figure 4.25)
Figure 4.24 shows resultant deterioration at the stage of (6) in the ﬂow chart above. This is
the expected worst case. Figure 4.25 shows the result of adjustment mentioned (7) in the ﬂow
chart. Note that the image is mostly recoverd to the original designed values (Figure 4.23).
This tolerance analysis shows that the WFC will qualify the requirement (D80 < 0.3”).

4.10

Performance Veriﬁcation Test

It would ideal to test the performance of the wide ﬁeld corrector using the exactly thesame incoming light from the telescope primary mirror but this is unrealistic. Instead, thefollowing two tests are considered[15].
Central image test
To run this test, an interferometer and a concave mirror should be needed. An interferometer
makes the incoming light beam into the measured lens system and forms a fringe pattern withreturning light beam. This light beam from the interferometer is adjusted to be focused at
the lens system focus position. A concave mirror reﬂects the light beam which comes throughthe measured lens system and gets the light beam to return to this lens system. And thisconcave mirror normal has to have the same angle with the light beam angle which come4.10. PERFORMANCE VERIFICATION TEST

Figure 4.23: Nominal design

Figure 4.24: Fabrication,assembling errors and environmental changes added

89

90

CHAPTER 4. WIDE FIELD CORRECTOR (WFC)

Figure 4.25: Adjust position/tilt of optics

through the lens system to make the return light beam focused at the lens system focus
position.
Furthermore a null lens should be needed for HSC test because Subaru telescope is RitcheyChretian system, then the primary mirror is a hyperboloid and has a spherical aberration.
Therefor even for a central image test, a null lens has to be made to correct this aberration.
The combined system of both the null lens and the measured lens system does not have theaberration. Actually there maybe a little residual aberration but it should be really small tomeasure precisely. If the measured lens system does not have any errors, such as fabrication
and assembly errors, this combined lens system has no aberration. But if the measured lenssystem has some errors this errors could be seen as a fringe pattern. If the measured lenssystem has lens decenter errors, coma aberration would be seen in a fringe pattern. Also if
the ﬁrst lens has a rotational symmetry surface error,spherical aberration would be seen ina fringe pattern. It is diﬃcult that ﬁx what causes an aberration exactly. But at least,we
can estimate what causes the aberration with some considerations as described above.
This test measure the central wavefront and is sensitive against front lenses tilt and decenter (See Figure 4.26).
The light from the interferometer pass through the null lens and WFC from the image
point,then reﬂected at the concave mirror and return to the interferometer. The null lens issmall compared to WFC lenses.
Total image test
It is very hard to implement the oﬀ axis test because it is also diﬃcult to design the nulllens for this test even if use many lenses and aspherical surface for the null lens and shiftconcave mirror.
Instead of the oﬀ-axis test, “total image” test is worked out (see Figure 4.27). This test
does not reconstruct the same light path with the original design. But this test is sensitive
against tilt and decenter of rear part of the optics elements because the beam of light isspread all over lens eﬀective diameter. Note that the distance to interferometer is long

4.11. COATING AND THE EXPECTED TRANSMISSION

91

Null Lens
Concave Mirror

Figure 4.26: Central image test

(about 1m), which means the measurement system would also be large.

Concave Mirror

Null Lens

Figure 4.27: Total image test

In fact, Canon did the the same tests for Suprime-Cam’s WFC. It is said that additional
adjustments and reassembly of WFC unit was not needed because the tests results was
acceptable. We hope this also happens in the HSC case.

4.11

Coating and the expected transmission

The expected transmission of AR-coat[16] designes are shown in Figure 4.28. Speciﬁcation
for the transmission is also shown in the Figure. There are six diﬀerent AR-coats dependon the lens material and curvature. Three combination of those designs are shown. It turns
out that AR-coat which qualiﬁes the requirements is feasible.
Figure 4.29 shows expected total transmission with coating error is calculated. Then we
can expect that the real coating will meet with the requirements at each ﬁeld angle andwavelength. Now we need to measure the actual transmission with test piece glass and ﬁnd
the best parameter, such as a gas pressure and so on, of vacuum evaporation coating device.

92

CHAPTER 4. WIDE FIELD CORRECTOR (WFC)

Transmission vs Wavelength
1
0.9
0.8
Transmission

0.7
0.6
0.5
0.4
Coating Design (a)0.3

Coating Design (b)0.2

Coating Design (c)
Specs

0.1
0
300

400

500

600

700
800
Wavelength (nm)

900

1000

1100

1200

Figure 4.28: Expected transmission curves of three diﬀerenct designs of the coating.

1.00
0.95

total
全系透過率
transmission

0.90
0.85
0.80
0.75
0.70
0.65
0.60
350

550

750

950

wavelength
波長 nm nm

1150

設計値
( 画角0
0)
画角(angle0)
designed
設計値( 画角00.375)
画角(angle0.375)
375)
designed
設計値( 画角00.525)
画角(angle0.525)
525)
designed
設計値( 画角00.675)
画角(angle0.675)
675)
designed
設計値
( 画角0
0.75)
画角(angle0.75)
75)
designed
最小値
( 画角0
0)
画角(angle0)
minimum
最小値( 画角00.375)
画角(angle0.375)
375)
minimum
最小値( 画角00.525)
画角(angle0.525)
525)
minimum
最小値( 画角00.675)
画角(angle0.675)
675)
minimum
最小値
( 画角0
0.75)
画角(angle0.75)75)
minimum
最大値
( 画角0
0)
画角(angle0)
maximum
最大値( 画角00.375)
画角(angle0.375)
375)
maximum
最大値( 画角00.525)
画角(angle0.525)
525)
maximum
最大値( 画角00.675)
画角(angle0.675)
675)
maximum
最大値
( 画角0
0.75)
画角(angle0.75)
75)
maximum
目標値
targetFigure 4.29: Result of error analysis of the coating4.12. SUMMARY

4.12

93

Summary

The major parameters of HSC WFC is summarized in Table 4.18 in comparison with thatof Suprime-Cam.

g
r
i
z
y

Suprime-Cam
0.5[deg]
450nm > 85%(∗1
546.1nm > 90%(∗1
850nm > 87%(∗1
0.18
0.09
0.11
0.15
0.17

HSC
1.5[deg]
400nm > 0.9%(∗2600nm > 89.0%(∗21050nm > 81.6%(∗2
0.16
0.15
0.15
0.15
0.20

g
r
i
z
y

0.21
0.13
0.13
0.16
0.19

0.20
0.19
0.19
0.19
0.22

ﬁeld of viewtransmission

Performance(D80”)
EL=90

Performance(D80”)
EL=30

focus length
ADC
Vignetting
Image curvature
Distortion
Lens weight
Ghost(illuminance ratio)

15000mm
18320mmlateral shift type ADC lateral shift type ADC
Non
max 25.6%
Non(plane image)Non(plane image)+0.9%+3.19%∼57kg
425kg< 1.1E-07< 5.4E-08(∗3

Table 4.18: Comparison of HSC WFC and SC WFC (*1:measured value (*2: designed value(*3:not include the ghost between ﬁlter surfaces

Bibliography[1] Canon Inc. 2008, HSC Report,HS40-045[2] Epps, H.W., Angel, J.R.P., Anderson, E 1984, in Very Large Telescopes, their Instrumentation and Programs, eds. M.-H. Ulrich, K. Kjär, Proc. IAU Coll. 79, p.519[3] Takeshi, K., Matsui, Y., Nariai, K., Ban, M. 1994, Japanese Laid Open Patent Application No. 6-230274[4] Nariai, K., Takeshi, K. 1994, SPIE, 2199, 532[5] Komiyama, Y. et al. 2004, SPIE, 5492, 525–532[6] Komiyama, Y. et al. 2006, SPIE, 6269, 106[7] Miyazaki, S. et al. 2006, SPIE, 6269, 9[8] Canon Inc. 2008, HSC Report, HS40-065[9] Canon Inc. 2007, HSC Report, No.20070928001[10] Canon Inc. 2009, HSC Report, HS40-120[11] Canon Inc. 2008, HSC Report, HS40-046[12] Canon Inc. 2008, HSC Report, HS40-081[13] Canon Inc. 2008, HSC Report, HS40-023[14] Canon Inc. 2008, HSC Report, Tolerance Analysis-080711[15] Canon Inc. 2008, HSC Report, HS40-100[16] Canon Inc. 2008, HSC Report, HS40-10295Chapter 5
HSC Camera Mechanics
5.1

CCD Dewar

The CCD dewar is the central component of HSC. It contains hundreds of CCDs which
receive photons emitted from astronomical objects. The CCD must be precisely arranged so
as not to spoil the excellent image quality realized by the wide-ﬁeld corrector. The CCDs
must be operated at a low temperature around –100◦ to be sensitive to faint objects by
reducing the dark current. The the CCD dewar must be properly designed taking various
aspects relevant to HSC into consideration.

5.1.1

Requirements

5.1.1.1

Optical Requirements

The design of the wide-ﬁeld corrector provides us an excellent image quality (D80 < 000 .22
from g 0 to Y bands) over the 1.5 degree ﬁeld of view at the telescope elevation between 30
and 90 degrees. The focal plane is ﬂat and its diameter is 495 mm. We have to install as
many CCDs as possible on the focal plane with minimum gaps between CCDs.
The F ratio of the incident light is 2.25 at the focal plane, which corresponds to the platescale of 0.893 mm/arcsec. The F ratio is rather fast and the focal depth is ∼ 34µm when 15µm pixel CCDs are adopted. This means that the CCD surfaces (i.e. focal plane) must be
ﬂat with an order of 10 µm. The total error budget assigned for the focal plane ﬂatness is
30µm (see Chapter 3). This budget is divided into the variation of CCD heights (25µm; see
Chapter 6) and the ﬂatness of the mounting plate of CCDs (18µm). The latter is relevantto the design of the CCD dewar. The ﬂatness error of the mounting plate is further divided
into the manufacturing error of the plate (15µm) and the deformation caused by the attitudeof the CCD dewar and the air pressure (10µm).
The back focal length of the wide-ﬁeld corrector is 190 mm. This space is shared by the
CCD dewar, the ﬁlter exchange system and the shutter assembly.
5.1.1.2

Mechanical Requirements

The prime focus unit, to which the CCD dewar and the wide-ﬁeld corrector are attached,is equipped with the hexapod system for the attitude control. This system is necessary
to correct the ﬂexure of the structures (including the telescope and the prime focus unit)
caused by the gravity and keep the optical axis of the prime focus unit aligned to that of the97

98

CHAPTER 5. HSC CAMERA MECHANICS

primary mirror at any telescope elevation. However, only one hexapod system is equipped
due to weight constraint. This means that, although the overall attitude can be corrected,
the relative position (tilt and displacement) between the wide-ﬁeld corrector and the CCD
dewar (CCDs) cannot be corrected.
We must design the CCD dewar taking the tilt and the displacement of the focal plane into
account. The error budget for the tilt (relative to the interface surface between the CCD
dewar and the dewar I/F frame) is assigned to be 10 arcsec at any elevation angle between 30
and 90 degrees (see Chapter 3). That for the displacement to the direction of the optical axis
(i.e., defocus) is 40µm. The error budget for the displacement to the direction perpendicular
to the optical axis is 15µm and the deviation of the displacement (as a function of the dewar
rotation angle at a ﬁxed elevation angle) must be less than 1µm.
The shutter assembly and a part of the ﬁlter exchanger system should co-rotate withthe dewar (CCDs) to ensure an accurate photometry. It is straight forward that thesecomponents are attached directly to the CCD dewar unless it is impossible. We must provide
appropriate interfaces to these components on the CCD dewar.

5.1.1.3

Thermal Requirements

The CCDs must be cooled down and operated at the temperature of –100◦ C to reducethe dark current generated in CCDs. Although the main component of the heat inﬂow to
the CCDs is the thermal radiation from the window (see Section 5.1.3), we have to makeevery eﬀort to reduce the thermal inﬂow, such as the thermal conduction and the thermal
convection, to the CCDs. In addition, the uniformity of the CCD temperature over the
focal plane is required since the temperature variation leads to the variation of the quantum
eﬃciency in particular at long wavelength around 1 µm. The variation must be less than
3 degrees across the focal plane. The stability of the CCD temperature when the ambient
temperature changes is another important issue.

5.1.1.4

Electrical Requirements

In order to securely handle hundred of CCDs, we plan to install the analog-to-digital conversion circuit inside the CCD dewar and reduce the number of cables and connections.
This also leads to a reliable operation since the electronics are protected inside the rigid
CCD dewar. The front-end electronics (FEE) is being designed as small as possible and,accordingly, the CCD dewar must be appropriately designed to accommodate large numbersof FEEs. A drawback to install the FEEs inside the CCD dewar is the heat generated fromthe FEEs. We have to introduce a mechanism to eﬃciently remove the heat from the FEEs
to the outside.

5.1.1.5

Weight and Space Constraint

The size of the CCD dewar, the shutter assembly and the ﬁlter exchange system determines
the size of the prime focus unit. Since the total weight is limited (3 tons), the outer sizeof the CCD dewar must be as small as possible while enclosing the 500 mm diameter focal
plane. The weight assigned to the CCD dewar is 200 kg.

5.1. CCD DEWAR

5.1.2

99

Mechanical Design

Figure 5.1 shows the current design of the CCD dewar. The dimension of the CCD dewar
is roughly 700 mm in diameter and 500 mm in height. Main components are labeled in the
exploded view (see Figure 5.1(c)).

(a)

Window
Assembly(b)

Wall
Assembly

Focal Plane
Assembly

FEE
Assembly

Back
Assembly(c)
Figure 5.1: (a): Overview of the CCD dewar. The dimension of the CCD dewar is roughly
700 mm in diameter and 500 mm in height. (b): Detailed view of the focal plane assembly.
(c): Exploded view of the CCD dewar. Main components of the CCD dewar are labeled.

5.1.2.1

Focal Plane Assembly

The focal plane assembly is the central component of the CCD dewar.

100

CHAPTER 5. HSC CAMERA MECHANICS

The base ﬂange is the base component of the CCD dewar to which other components
are assembled. It also functions as the interface to the dewar frame which is provided by
Mitsubishi Electric Co. It has to be stiﬀ enough to prevent the entire dewar structure fromﬂexure. The material of this ﬂange is planned to be an aluminum of 32 mm thick.
The cold plate is the mounting plate to which 116 CCD packages are attached. Its dimension is 620 mm in diameter (with two side cuts) and 20 mm in thickness. We choose
silicon-carbide (SiC) SC1000 for the material of the cold plate taking the thermal conductivity, the Young modulus, the thermal expansion and the manufacturability into consideration.
SiC is a suitable material for the cold plate since it has high thermal conductivity (200 W
m−1 K−1 ) and high Young modulus (440 GPa). Kyocera Co. is capable to manufacture SiC
plate of this size with the ﬂatness of 15 µm. Since the CCD package is made of aluminumnitride (AlN), the material of the cold plate must have CTE close to that of AlN. The CTEof AlN (4.6 ppm) is diﬀerent from that of SiC (3.7 ppm), we have to carefully design theholes on the SiC cold plate where CCD packages are screwed so that CCD packages are
not damaged when the cold plate is cooled down to –100◦ . Our calculation shows that thediﬀerence of the heat shrink between AlN and SiC for the length between two positioning
pins on CCD package is comparable to the heat shrink of the diameter of Titanium positioning pins on CCD packages. We conclude that SiC can be used for the material of thecold plate. Note that the maximum available size of AlN plate is 500 mm in diameter as far
as our inquiry went. The cold plate is covered by radiation shields which are made of highemissivity material to reduce the heat inﬂow to the cold plate.
The dimension of CCD package is 31.54 mm × 66.825 mm with 9 mm thick. Threepositioning pins with screw, which are made of Titanium, are glued to the AlN base of the
package to determine the position on the cold plate and rigidly ﬁx to the cold plate. We set
the gap between adjacent CCD packages to be 0.2 mm taking the handling of CCD packages
into account. The gap in case of Suprime-Cam is 0.3 mm and we did not have any diﬃcultyin handling CCD packages. PGA of 59 pins is equipped on the back side of the CCD package
and a FPC cable is connected to the PGA, which makes the CCD package 4 side buttable.
To feed FPC cables, feed-through holes are required to the cold plate in addition to thepositioning holes of CCD packages. The length of FPC cable must be ∼ 300 mm for easy
installation (see Section 5.1.7).
The support posts are the parts which connect the cold plate to the base ﬂange. The
selection of this part is very important since it has to be a stiﬀ material to keep the focal
plane from deformation, while minimizing the thermal inﬂow by conduction from the baseﬂange at the room temperature to the cold plate at –100◦ C. Table 5.1 summarizes the physicalproperties of candidate materials for the support post. When a simple shape rectangular solid
is considered, the thermal conduction is proportional to the area while the ﬂexure is inverse
proportional to the area. This means that a material which has higher ratio of the Youngmodulus to the thermal conductivity is more appropriate for the material of the support
post. But, when a no simple shape solid is considered, this is not the case. We ﬁnd thatthe stainless steel tube can reduce the thermal inﬂow by conduction compared with simple
shape solids while maintaining the stiﬀness. We plan to employ truss structure to support
the cold plate as shown in Figure 5.1 (b). Each bar is made of 160 mm long stainless steel
tube of 21 mm outer diameter with 0.5 mm radial thickness. The structural and thermal
performance of this stainless steel tube truss structure is summarized in Section 5.1.2.7 and5.1.3, respectively.

5.1. CCD DEWAR

101

Material Young Modulus Thermal Conductivity
CTE
(GPa)
(W/m K)(10−6 /K)PC
2.3
0.12-0.19
67.5
G10
21-24
0.38-0.47
10-20
SUS304
193
16.2
17.5
Zirconia
250
1.9-4.0
10.8
ArFRP
70
0.23
-1.5
CFRP
150-420
4-5-1
Table 5.1: Physical properties of candidate material for the support post.
5.1.2.2

Wall Assembly

The wall assembly is in short a support structure of the window assembly. It is equippedwith feed-through connectors for the FEEs inside the CCD dewar. The number of connectors
required is estimated to be 24 for the CCD readout and several for the house keeping. We
make a room to attach 30 connectors on the wall assembly.
The wall assembly also works as an interface to the shutter assembly and a part of theﬁlter exchange system. We plan to attach these components with M5 screws. The detailed
design of the interface is being discussed and will soon be ﬁxed.
5.1.2.3

Window Assembly

The window assembly consists of a glass window and the support structure of the window.
We plan to employ tapered shape window as shown in Figure 5.2 (a). This shaped window
has a merit that the window can be placed close to the CCD surface and the back focal
space can be assigned for the shutter assembly and the ﬁlter exchange system. We ﬁrst try
to test the tapered shape window as a prototype.
The tapered surface is made with an concentric tolerance of 0.1 mm. The surface issupported by two PTFE rings and they determine the tilt of the window relative to thewindow support. An O-ring is embedded between these PTFE rings to work as a vacuum
seal. We have to be careful that the window must be well aligned to the support structure
when we attach the window to the support structure. Otherwise, we may have a vacuum
leak or a damage on the window in worst case. To avoid such problems, we plan to (1)measure the three dimensional shape of the tapered surface and the support structure, and(2) adjust the height of the PTFE rings with shims to ensure the rings to contact properly tothe tapered surface of the window, and (3) push the window down to the support structure.
The window is made of synthetic silica with anti-reﬂection coating on both side. The
larger and smaller diameters are 586 mm and 550 mm, respectively, and the thickness is 37
mm (clear aperture 537 mm, supported at 556 mm). The support structure which is madeof aluminum holds the window from the periphery. The support structure does not help toreduce the ﬂexure of the window when the CCD dewar is evacuated since the window is notglued (or rigidly attached) to the structure and is supported by O-ring and PTFE rings.
The thickness of the window is determined as follows. A window of radius R and thickness

102

CHAPTER 5. HSC CAMERA MECHANICSt which is supported at its circumference suﬀers the stress described as,
3P (3 + ν)σmax =
8

µ ¶2
R
,
t

(5.1)

where P is the pressure diﬀerence between inner and outer surfaces and ν is the Poisson
ratio. The sag at the center of the window l is calculated as,l=3P (1 − ν)(5 + ν) R4
,
16E
t3

(5.2)

where E is the Young modulus. We carried out the calculation for diﬀerent thickness andfound that the safety factor of 7.2, which is the adopted value for Suprime-Cam, is realized
when we assign 37 mm for the window thickness. We therefore choose 37 mm for the
thickness. In this case, the sag at the center of the window amounts to 138.1 µm and,hence, the window deforms to a meniscus shape with a radius of 280 m. We conﬁrmedthat the ﬂexure of the window gives no signiﬁcant eﬀect on the image quality. The adopted
properties of silica are summarized in Table 5.2. We also carried out the structural analysisof the window in Section 5.1.2.7.
A mechanism to brow dry air is attached to the support structure to keep the window dry
and prevent the window from dew condensation (not shown in Figures 5.1 and 5.2). For
this purpose, dry nitrogen generated at the telescope compressor room will be continuously
supplied to the prime focus through the pipe.
The optical speciﬁcations of the window are as follows. The transmitted wavefront error
must be less than 100 nm for 30 mm aperture at any position over the ﬁeld of view. This
speciﬁcation is same as that for ﬁlters (see Chapter 7). The parallelism between inside andoutside surfaces must be less than 40 arcsec. This requirement comes from the image quality
error budget. Anti-reﬂection coating is applied on both side of the window. The expectedtransmission is 97-98 % for the wavelength between 400 nm and 1000 nm. Figure 5.2 (b)shows an actual reﬂectance measurement, provided from one of coating vendors, of candidate
anti-reﬂection coating. This coating has already be applied for larger glasses than 500 mm
and shows good performance.

Breaking strength
4.892 ×107 Pa
Young modulus (E) 6.969 ×1010 Pa
Poisson ratio (ν)
0.17
Table 5.2: Adopted physical properties of silica.

5.1.2.4

FEE Assembly

The FEE assembly holds the front-end electronics inside the CCD dewar eﬃciently andrigidly. It also works as an heat exchanger to remove the heat generated inside the CCD
dewar.

5.1. CCD DEWAR

103

(a)

(b)

Figure 5.2: (a) Cross sectional view of the window assembly. (b) An actual reﬂectancemeasurement of planned anti-reﬂection coating.
Figure 5.3 (a) shows the overview of FEE assembly. There are two types of FEE boards:
One handles 4 CCDs (board dimension: 240 mm × 150 mm × 2 mm) and the other handles
6 CCDs (360 mm × 150 mm × 2 mm).
FPC cables from CCD packages are connected to 65 pin micro D-sub connectors on thebottom side of the FEE board. The FEE board has one 9 pin MDM connector one end of
the board which is used as a signal line. On the other end, a 100 pin connector is attachedand connected to the bus connector which is installed on the part called FEE plate and the
DC power is supplied from the bus (Figure 5.3 (b)). A support bar which is made of copperis attached at the top end of the FEE board to ﬁx it to the FEE plate by two M3 screws.
The bar also works to transfer the heat generated from the FEE board to outside eﬃciently.
For more details on the FEE, readers should refer to Chapter 6.
The FEE plate holds FEE boards and bus boards rigidly and transfer the heat generated
on FEE boards to the outside of the CCD dewar. It is made of A6061 aluminum alloy
which has higher thermal conductivity than ordinary aluminum alloy. The details on theheat transfer from FEE boards to the outside is summarized in Section 5.1.3. As shown in
Figure 5.3 (b), a part with insertion slots is attached at the other end of the bus board.
These slots are necessary to constrain the movement of FEE boards (with a help of two M3screws on the support bar) and ensure the stable connection between FEE boards and thebus board.
In the FEE assembly, there are two rooms reserved for the cold path between the coldhead and the cold plate. The size of the each room is about 500 mm × 65 mm. The thermal
analysis (see Section 5.1.3) shows that wide contact area on the cold plate is necessary tokeep CCD temperature uniform over the focal plane.
How FEE boards are attached to the FEE plate is summarized as follows. (1) First, 4
or 6 FPC cables from CCDs are connected to the connectors on a FEE board. Note that
the FPC cables from CCDs must be ∼300 mm long. (2) While folding the FPC cables, theFEE board is aligned to the insertion slot (Figure 5.3 (b)). (3) The FEE board is slided andconnected to the bus board. At this time, the FEE board is inserted in the slot. (4) Two
M3 screws on the support bar is attached and the FEE board is fully ﬁxed. (5) This process

104

CHAPTER 5. HSC CAMERA MECHANICS

is repeated and, ﬁnally, all the FEE boards are ﬁxed to the FEE plate (see Figure 5.3 (c)).

(a)

(b)

(c)
Figure 5.3: (a) Overview of the FEE assembly without FEE boards. (b) Detailed viewaround the bus board. (c) Schematic view which shows how FPC cables from CCDs areconnected to FEE boards. Note that 9 FPC cables for the central FEE group are drawn inthe ﬁgure.

5.1.2.5

Back Assembly

The back assembly is where coolers and vacuum maintenance instruments are attached. An
ordinary vacuum valve of ISO-KF 40 size (e.g., Pfeiﬀer EVB 040 SA) and a full range vacuum
gauge (e.g., Pfeiﬀer PKR 251) are attached. We also plan to adopt a noble pump Varian
VacIon Plus 20 which has a pumping speed of 20 l/s. Details of the cooler and the pump aresummarized in Section 5.1.4 and 5.1.5. On the back assembly, a rack which holds controllers
and power supplies (such as BEEs and CCD power supplies) is also attached.

5.1. CCD DEWAR
5.1.2.6

105

Focal Plane Layout

(a)

(b)

Figure 5.4: (a): The layout of CCDs on the cold plate. (b): The eﬀective area of the focal
plane.
We plan to install 116 CCDs to cover the 1.5 degree ﬁeld of view focal plane. Figure 5.4(a)shows the layout of 116 CCDs on the focal plane. The inner dashed circle represents the
1.5 degree ﬁeld of view whose physical size is 495 mm. As shown in the ﬁgure, 116 CCDs
perfectly cover 1.5 degree ﬁeld of view. The outer dashed circle represents the area where
CCDs receive photons. The diameter is 530 mm in physical and 1.6 degree ﬁeld of view.
As explained in Chapter 4, the vignetting gets worse as the radius increases and it is 25%
at 0.75 degree (i.e., 1.5 degree diameter), then steeply degraded to ∼40% at the 0.8 degreeedge. The solid line circle represents the outer diameter of the CCD dewar (700 mm). The
set of 116 CCDs is not allocated on the right center of the cold plate; 2 mm oﬀset to thebottom dirrection in the ﬁgure. This allocation allows the object on the optical axis fallingon the image area of CCD. Otherwise, it is fallen on the gap between the CCDs.
The dimension of a CCD package is 31.54 mm × 66.825 mm with 9 mm thick and thegap between CCD packages is planned to be 0.2 mm. Since the light-sensitive area of CCD
is 30.72 mm × 61.44 mm, the eﬀective gaps between CCDs are 1.02 mm (11 arcsec) in 2k
direction and 5.585 mm (62 arcsec) in 4k direction, respectively. Figure 5.4(b) shows theligh-sensitive area (white) and the gap (black) of the focal plane. The ﬁlling factor of the
focal plane is calculated to be about 90 %.
We plan to use some CCDs at the edge of the focal plane for special purpose; 4 CCDs
for auto-guiding (AG) and 8 CCDs for auto-focusing (AF). The AG CCDs are controlled bydedicated FEEs which are tuned to readout very quickly (see Chapter 6). The AF CCDs havediﬀerent height from normal CCDs; 4 CCDs are low and other 4 CCDs are high. From thedefocused images obtained by AF CCDs, we can estimate the best focus position. Although

106

CHAPTER 5. HSC CAMERA MECHANICS

these CCDs cannot be used for the science unfortunately, only small portion of 1.5 degreeﬁeld of view is occupied by these CCDs. We believe there is no signiﬁcant eﬀect for thescience. More details on auto-guiding and auto-focusing are summarized in Chapter 8.
5.1.2.7

Structural Analysis

We carried out the structural analysis of the CCD dewar using the ﬁnite element analysis(FEA). What we have investigated are the deformation of the cold plate, the base ﬂange
and the window. Table 5.3 summarizes the mechanical properties used for the FEA.

Material
Aluminum
StainlessCopperSiC
AlN
Fused Silica

Class

Density Young Modulus Poisson Ratio(kg/m3 )(GPa)
A5052-H34
2700
70.3
0.33
A6061-T6
2700
68.9
0.33
SUS304
8000
193
0.29
C10100
8940
115
0.31
SC1000
3160
440
0.17
AN216A
3400
320
0.24
2200
70
0.17

Allowable Stress
(MPa)
214
276
215
69
450
310
48

Table 5.3: Adopted mechanical properties used for the structural analysis.

The Flexure of the Cold Plate
We carried out the deformation analysis to see if the deformation of the cold plate satisﬁes
the requirements which come from the image quality error budget. FEAs were performed
for elevation angles (El) of 0, 30, 60 and 90 degrees and for diﬀerent rotator angles (Rot) of
0, 30, 60, 90, 120, 150 and 180 degrees, where the rotator angle (Rot) is the angle between
y axis of the telescope coordinate system and y 0 axis which is the local dewar coordinate
system (see Figure 5.5). The analysis for diﬀerent rotator angles are needed since we employ
the truss structure for support posts and the deformation is likely to be dependent on the
direction of the gravity. Note that we do not use the result of El=0 for the judgment of ourdesign and use it for reference.
The conditions of the FEA are summarized as follows.
1. The back side of the base ﬂange is fully constrained.
2. The whole assembly (including window, wall, back assemblies) is included in the analysis.
3. The perfect contact is assumed between parts.
4. The ambient air pressure is 0.6 atm (demonstrates the summit condition).

5.1. CCD DEWAR

107

Figure 5.5 shows the displacement for z axis (∆z 0 ). The resultant deformation is not
simple; the cold plate is pushed or pulled locally by six support posts and warped. We
obtain the ﬂatness, tilt and displacement as follows: (1) First, determine the best-ﬁt plane(z 0 = ax0 + by 0 + c) from the ﬁtting of the data set (x0 , y 0 , ∆z 0 ). (2) We deﬁne the ﬂatness ofthe cold plate surface as the deviation from the best-ﬁt plane (i.e., |maximum deviation| +|minimum deviation|). (3) We deﬁne the tilt
√ of the cold plate surface as the tilt between thebest-ﬁt plane and the x0 − y 0 plane (i.e., a2 + b2 radian). (4) We deﬁne the displacementof the cold plate surface for y (∆y) as the displacement of the center of the cold plate fory direction. (5) We deﬁne the displacement of the cold plate surface for z (∆z) as thedisplacement of the center of the cold plate for z direction.
Table 5.4 summarizes the best-ﬁt plane and ﬂatness, tilt and displacement calculated from
FEA. Table 5.5 is the summary table of the FEA in addition to the requirements. The worst
ﬂatness is 6.5 µm at El=30 and Rot=180, which is less than the assigned error of 10 µm.
The maximum tilt is 0.83 arcsec at El=30 and Rot=120, which is far less than the assigned
error of 10 arcsec. The maximum displacement for z direction is 7.2 µm at El=90, which isalso less than the assigned error of 40 µm. The maximum deviation of the displacement fory axis is 0.7 µm at El=30, which satisﬁes the assigned error of 1 µm. We therefore conclude
that the ﬂexure of the CCD dewar is acceptable.

El
90
60

30

0

Rot
0
0
30
60
90
120
150
180
0
30
60
90
120
150
180
0
30
60
90
120
150
180

a
-4.94E-09-2.69E-08-1.15E-06-1.97E-06-2.26E-06-1.95E-06-1.11E-06
1.66E-08-4.30E-08-2.06E-06-3.52E-06-3.91E-06-3.48E-06-1.99E-06
3.22E-08-4.92E-08-2.30E-06-3.93E-06-4.52E-06-3.89E-06-2.22E-06
3.72E-08b
1.02E-07-2.10E-06-1.80E-06-9.91E-07
1.01E-07
1.19E-06
1.98E-06
2.26E-06-3.75E-06-3.34E-06-1.90E-06
6.03E-08
2.00E-06
3.41E-06
3.80E-06-4.41E-06-3.80E-06-2.19E-06-9.30E-09
2.16E-06
3.75E-06
4.31E-06c-0.006338-0.005671-0.005699-0.005770-0.005867-0.005963-0.006032-0.006056-0.004234-0.004086-0.004214-0.004574-0.004559-0.004683-0.004902-0.000002-0.002468-0.002611-0.002805-0.002996-0.003135-0.003183

Flatness
0.002179
0.002346
0.003054
0.003549
0.003794
0.003873
0.0041170.004182
0.004963
0.005835
0.006278
0.006392
0.006475
0.006301
0.005935
0.000006
0.006795
0.007069
0.007364
0.007121
0.006695
0.006100Tilt
0.021000
0.432672
0.440032
0.454699
0.466707
0.470488
0.468293
0.466291
0.773012
0.808930
0.824824
0.806912
0.827683
0.814667
0.784029
0.908932
0.916581
0.929047
0.931606
0.918315
0.898474
0.888998∆x0.000000
0.000000-0.000223-0.000328
0.000260-0.000145-0.000048
0.000000
0.000000-0.000225-0.000274
0.0001300.000061
0.0001050.000000-0.000003-0.000172-0.000154
0.000400
0.000223
0.000212
0.000000∆y0.000310-0.003150-0.003235-0.003448-0.003630-0.003709-0.003702-0.003690-0.005780-0.006111-0.006354-0.006300-0.006466-0.006319-0.006060-0.006880-0.006982-0.007186-0.007280-0.007146-0.006912-0.006800∆z-0.007180-0.006120-0.006180-0.006350-0.006580-0.006820-0.006990
0.007050-0.004150-0.004000-0.004300-0.004950-0.005130-0.005440-0.005760-0.001800-0.001920-0.002260-0.002730-0.003190-0.003530-0.003660

Table 5.4: The result of FEA. The unit of Tilt is arcsec and ∆x, ∆y, ∆z is mm.

The Flexure of the CCD Dewar System
We carried out the deformation analysis to see if the ﬂexure caused by the air pressure
at 1 atm does not damage the CCD dewar. The conditions of the FEA are summarized asfollows.
1. The whole system is freely placed (no constraint).

108

CHAPTER 5. HSC CAMERA MECHANICS

(a)

(b)

Figure 5.5: (a): The displacement for z direction when the elevation angle of 30 degrees andthe rotator angle of 0 degree. (b): Same ﬁgure, but for the case of the rotator angle of 90
degrees.

El
90
60
30
0
Req.

Flatness
Tilt h∆xi R(h∆xi) h∆yi R(h∆yi) h∆zi R(h∆zi)[µm] [arcsec]
[µm]
[µm]
[µm]
[µm]
[µm]
[µm]
2.179
0.021 0.000- 0.310- -7.180
4.182
0.470 -0.069
0.588 -3.509
0.559 -6.584
0.930
6.475
0.828 -0.029
0.404 -6.199
0.686 -4.819
1.760
7.364
0.932 0.072
0.572 -7.027
0.480 -2.727
1.860
10
10
15
1
15
1
40
-

Table 5.5: The summary of the deformation analysis. R(h∆xi) represents the range of h∆xi.
The requirements are also summarized in the bottom row.
2. The whole assembly is included in the analysis.
3. The perfect contact is assumed between parts.
4. The ambient air pressure is 1 atm (demonstrates Mitaka or Hilo condition).
Table 5.6 summarizes the maximum stress, safety factor and the deformation calculated
from FEA. The base ﬂange is deformed by 41.8 µm because the ﬂexure of the back assemblyis large and transfered to the base ﬂange. The back assembly should be re-designed so thatthe ﬂexure of the base ﬂange is minimized. The most fragile part of the CCD dewar is thewindow. The largest stress occurred at the center of the window and it amounts to 6.6

5.1. CCD DEWAR

109

MPa. The safety factor is 7.2 which is consistent with the analytical calculation done in
Section 5.1.2.3. We do not ﬁnd any problem for the cold plate. We therefore conclude thatour design is acceptable.

Base Flange
Cold Plate
Window

Maximum Stress [Pa] Safety Factor
4.5E+661.3
2.8E+6155
6.6E+6
7.2

Deformation [µm]
41.8
151

Table 5.6: The result of FEA for 1 atm case.

5.1.3

Thermal Design

5.1.3.1

Analytical Calculation

The heat inﬂow by the thermal radiation is calculated as
Q̇ = σA(TH4 − TL4 )²H ²L²H + ²L − ²H ²L(5.3)

where σ is the Stephan-Boltzmann Constant (5.67 × 10−8 W m2 K4 ), A is the area ofthe object, ² is the emissivity, T is the temperature, and the suﬃx H, L represents thehigher/lower temperatures. In our case, CCDs and the cold plate are cooled down to –100◦ C(TL = 173 K) and the surrounding structures such as the wall assembly and the window
assembly are all at the summit temperature (10◦ C, TH = 283 K) except for the FEE boards(50◦ C, TH = 323 K) .
The heat inﬂow by the thermal radiation can be signiﬁcantly reduced if we can reduce theemissivity (² << 1). The metals are basically low emissivity materials but the cold plate(ceramics), CCDs (silicon) and the window (silica) are not. The cold plate can be covered byradiation shields to reduce the heat inﬂow, but this is not the case for CCDs and the window.
Therefore, the main component of the heat inﬂow by the thermal radiation is occurred
between CCDs and the window. The emissivity of the CCD surface is measured to be 0.5,
from our measurement based on the calorimetric experiment and from the infrared reﬂectancemeasurement carried out at JAXA[3]. The heat inﬂow amounts to 38.3 W assuming that theemissivity of the silica window is unity (Table 5.7). The other components are also calculated
and found to be smaller than the radiation between CCDs and the window (Table 5.7). Notethat aluminum plate (emissivity ²=0.06) is adopted for the radiation shield on the cold platein this analysis.
The heat inﬂow by the thermal conduction is calculated as
Z
A
A TH
(5.4)
κdT ' κ̄(TH − TL )
Q̇ =
L TL
L
where κ is the thermal conductivity of the object, A is the area of the object, L is thelength of the object, T is the temperature, and the suﬃx H, L represents the higher/lower
temperatures. We have to consider support posts and FPC cables which connect between

110

CHAPTER 5. HSC CAMERA MECHANICS

lower temperature part (cold plate/CCD) and higher temperature part (base ﬂange/FEE).
As summarized in Table 5.7, the heat inﬂow by the thermal conduction is 1/10 lower thanthat by the radiation, but it cannot be negligible. The heat inﬂow through support posts is
reduced to 2.15 W from 7.56 W, which is estimated for the conceptual dewar design[2].
Another heat source is the heat generated from CCDs. This is calculated to be 0.024 W
for each CCD and, hence, total heat generation amounts to 2.78 W.
In conclusion, total heat inﬂow to the cold plate is estimated to be 52.25 W. The cooler
must be capable to exhaust this amount of heat.

SourceWindow – CCD
FEE – Cold Plate
Wall – Cold Plate

A (m2 )
0.244?
0.135
0.263??

²H²L Q̇ (W)
1
0.5
38.3
1
0.06
4.57
0.06 0.06
2.55(a) Thermal Radiation
SourceA (m2 )
Support Posts 1.93×10−4†
FPC Cables
1.23×10−5‡

L (m)
0.160
0.3

κ (W/m K) Q̇ (W)
16.2
2.15
420
1.90

(b) Thermal Conduction
SourceQ̇ (W)
CCD
116×0.024 = 2.78(c) Heat Generation
Table 5.7: (a): Heat inﬂows by the thermal radiation. (b): Heat inﬂows by the thermal
conduction. (c): Heat generated at cold side.
Note: ?: Total area of CCD surface: 116×0.03154×0.066825 (m2 ). ??: The total area ofthe cold surface excluding the CCD surface and the area facing to FEE boards. In thiscalculation, the cold plate is assumed to be 620 mm in diameter and 20 mm thick. †: 6
stainless steel tube of 21 mm outer diameter and 0.5 mm thick. ‡: FPC cable consists of 59
copper wires. The cross section of a wire is 100µm × 18µm. 116 FPC cables are connected
between CCDs and FEE boards.

5.1.3.2

Simulation

We carried out the thermal analysis of the CCD dewar using FEA. What we have investigated
are the CCD temperature, the thermal stress analysis of the window and the temperatureof the FEE assembly. Table 5.8 summarizes the thermal properties adopted for the FEA.
CCD Temperature
Figure 5.6 (a) summarizes the condition of the FEA. In this model, the temperature of
the base ﬂange is ﬁxed to 283 K. The heat coming from FEE boards through FPC cables

5.1. CCD DEWAR

Material

111

Class

CTE Thermal Conductivity(ppm)(W / m K)
Aluminum A5052-H34 23.8
138
A6061-T6
23.6
167
Stainless
SUS30417.3
16.2
Copper
C1010017383
SiCSC10003.7
200
AlN
AN216A
4.6
150
Fused Silica
0.4
1.4
FEE
CCD

Emissivity
0.06
0.06
0.074
0.02
1
1
1
1
0.5

Table 5.8: Adopted thermal properties used for the thermal analysis.
is treated as a heat generation in CCDs and the total heat generation of CCDs amount to4.68 W. The cold plate receives the heat inﬂow by radiation from wall assembly and FEE
boards and that by conduction through support posts. The window, which faces the cold
CCD surface, is cooled by the radiation cooling and the temperature of the window center
reaches to 257 K. We found that 43.68 W heat exhaust by coolers is required to keep the
CCD surface 173 K.
In the steady state, the temperature of CCDs ranges from –100.0◦ C to –101.0◦ C over the
focal plane as shown in Figure 5.6 (b). The variation of CCDs is only 1 K and satisﬁes therequirement that the variation of the CCD temperature must be less than 3 degrees acrossthe focal plane. We conclude that our design is acceptable.
In reality, CCD temperature will be slightly higher since the perfect thermal contact between CCDs and the cold plate is assumed in this model. The diﬀerence of temperature
measured from Suprime-Cam is 2∼3◦ C.
Thermal Stress Analysis of the Window
As explained in the previous paragraph, the window is cooled by the radiation coolingand the temperature of the window drops below the ambient temperature. The window
will suﬀers a stress caused by the thermal variation in the window. We calculated thetemperature distribution in the window and the stress caused by the thermal variation.
Figure 5.7 (a) summarizes the condition of the FEA. In this model, the temperature of
the tapered surface is ﬁxed to 300 K. The outside surface of the window receives a heat fromambient air by heat transfer with a rate of 4.6 N m−1 s−1 K−1 . The inside surface of thewindow emits radiation to the CCD surface and cooled. Figure 5.7 (b) shows the ultimate
temperature distribution of the window. The temperature at the center of the outside surface
is 21 degrees lower than the ambient temperature. This indicates that the window surface is
so cold to have a dew condensation. To avoid the dew condensation, we have to blow dry air
to the window surface continuously. The temperature at the center of the inside surface is
26 degrees lower than the ambient temperature. This temperature drop is welcome from thethermal point of view (i.e. heat inﬂow by radiation is reduced). We calculated the thermal

112

CHAPTER 5. HSC CAMERA MECHANICS

(a)

(b)

Figure 5.6: (a) The analysis condition of the model. (b) The result of thermal analysis.
stress caused by the temperature gradient and it amounts to 1.3×105 Pa, which is lower
than the stress caused by the air pressure (6.6×106 Pa, see Section 5.1.2.7). We therefore
conclude that our design is acceptable.

(a)

(b)

Figure 5.7: (a) The analysis condition of the model. (b) The temperature distribution of
the window.

5.1. CCD DEWAR

113

Temperature Distribution of the FEE Assembly
We plan to exhaust the heat generated by FEE boards inside the CCD dewar to the outsideby thermal conduction. The key element is the FEE plate and the support bar on the FEE
board. We have carried out the FEA analysis of the FEE assembly.
Figure 5.8 (a) summarizes the condition of the FEA. In this model, the temperature of
the periphery of the FEE plate is ﬁxed to 300 K. the support bar receives 1 W per CCDfrom the FEE board as calculated in Chapter 6. The FEE plate is made of aluminum alloy(A6061) and the support bar is made of copper. Figure 5.8 (b) shows the result. The
maximum temperature of the support bar is 41.1 degree C, which satisﬁes the requirement
coming from the FEE design (see Chapter 6). We therefore conclude that our design of FEE
assembly is acceptable.

(a)

(b)

Figure 5.8: (a) The analysis condition of the model. (b) The temperature distribution of
the FEE assembly (FEE plate and support bars).

5.1.4

Cooler

The result of thermal analysis indicates that we have to have a cooler which is capable
to absorb ∼53 W heat at the temperature of –100◦ C. However, only limited type of such
powerful coolers are available.
One choice is the GM type coolers. This type of cooler is world-wide prevailing and employed by many instruments at many observatories. One demerit is that the cooler requiresa large compressor which must be placed on the ground. This means that very long He tubes
between the prime focus and the Cassegrain ﬂoor are necessary and an appropriate routing
of He tubes and suﬃcient cable wrapper must be prepared.
The other choice is the large pulse-tube type coolers which have a compressor directly
connected to the cold head. This type of coolers has merits of long life time and low
vibration in contrast to the GM type coolers. We surveyed this type of coolers and foundthat Fuji Electric Systems had developed a large pulse-tube cooler which has the cooling

114

CHAPTER 5. HSC CAMERA MECHANICS

power of 8 W at 77 K. We have tested the cooler and found that the cooler has the coolingpower of 41 W at –100◦ C. According to the engineer of Fuji Electric Systems, the cooler
was optimized to the use at 77K and will have the cooling power of 48 W at –100◦ C after
the optimization to the use at –100◦ C. Then, 2 coolers are enough to cool the CCDs on the
focal plane of HSC suﬃciently. In this case, the stability of the CCD temperature when theambient temperature is changed can be guaranteed since the cooling power is almost twice
more than the heat load. Since this type of coolers have a merit that we do not need long
He tubes, we determined to employ large pulse-tube coolers and pursue the realization ofthis type of coolers.
The prototyping work is in progress at Fuji Electric Systems. The prototype is now
assembled and is being tested (Figure 5.9). We will have the optimized cooler by the end of
March 2009.

Figure 5.9: The prototype cooler in test.

5.1.5

Vacuum Maintenance

When Suprime-Cam was ﬁrst attached to the prime focus of Subaru Telescope in 1999, cooled
charcoal was glued to the cold surface inside the dewar to keep the vacuum. However, it did
not work as we expected because the cooler of Suprime-Cam was not so powerful to cool the
charcoal work. We abandoned charcoal and instead employed an active vacuum pump, an
ion pump.
Suprime-Cam is now equipped with an Noble pump from Anelva Co. which have a exhaust
power of 20 l/s. The Noble pump works quite well to keep the vacuum down to 10−7 Torrwhen the cooler is working. The vacuum is kept so high that we do not have to care about
the life time of the pump (the life time is estimated to be more than 200000 hours). We
have had no particular trouble with this Noble pump for these 8 years.
We plan to attach 20 l/s ion pump to HSC as well. The attainable vacuum is estimated tobe 10−6 Torr, assuming that the amount of out gas is proportional to the number of CCDs
and HSC will have 10 times more out gas than Suprime-Cam (10 CCDs to 100 CCDs). The5.1. CCD DEWAR

115

life time is ∼20000 hours in this case, and 2000 hours even if the attained vacuum is 10−5
Torr.
The most compact and light one is VacIon plus 20 from Varian Co. which has a pumping
speed of 20 l/s for nitrogen (Figure 5.10). The pump and the full controller weight 7 kg
and 10 kg, respectively. The compact controller weights only 2.3 kg though some functions
such as serial communication are omitted. We would like to use full controller if the weightbudget permits, but the compact one is acceptable.

(a)

(b)

Figure 5.10: (a) VacIon plus 20 from Varian Co. (b) Pumping speed for nitrogen plottedagainst pressure.

5.1.6

Metrology of the Focal Plane Flatness

The surface of the cold plate is measured by contact type coordinate measuring machine(CMM) LEGEX 9100 by Mitsutoyo Co. (see Figure 5.11(a)). This CMM has 1m measurable
working area and the measurement error is 0.35+L/1000 µm, where L is the size of object.
This means the surface of the cold plate (620 mm in diameter) can be measured with 1 µm
accuracy. This CMM is available in advanced technology center (ATC) of NAOJ. Note thatthe cold plate must be measured as an assembly with the base ﬂange and the support posts
since the surface shape of the cold plate would be distorted when assembled.
We also have a non-contact type CMM in ATC, NH-3SN from Mitaka Koki Co. (seeFigure 5.11(b)). This CMM has 30 cm measurable working area and attains the measurement
accuracy of 1 µm. We plan to measure the surface shape of all CCD packages and the height
relative to the back side of the package. Based on our experience on the CCD upgrade ofSuprime-Cam held in 2008 summer, this CMM is capable to measure the surface of the CCD
package quite well.
Taking the surface shape of the cold plate and the height (and the surface shape) dataof CCD packages into account, we will try to determine the optimal CCD allocation on thecold plate to make the focal plane as ﬂat as possible.

116

CHAPTER 5. HSC CAMERA MECHANICS

(a)

(b)

Figure 5.11: (a) Legex 9100 from Mitsutoyo Co. (b) NH-3SN from Mitaka Koki Co.

5.1.7

Assembly Sequence

Figures 5.12 and 5.13 summarize the assembly sequence of the CCD dewar.
CCD Installation
At ﬁrst, the base ﬂange is prepared and ﬁxed to the coordinate measuring machine (Figure 5.12(a)). Support posts and the cold plate are assembled on the base ﬂange and adjusted
so that the cold plate surface is parallel to the base ﬂange (Figure 5.12(b)). This assemblyis attached to a dedicated dolly which has tumbling mechanism (not shown in the ﬁgure).
At this time, radiation shields are covered on the cold plate. A CCD package with insertion
pins is installed on the cold plate and is screwed from bottom side of the cold plate withtorque wrench (Figure 5.12(c)). This process is repeated for 116 times and the focal plane
is completed (Figure 5.12(d)). A CCD protection lid (not shown) is covered for safety. Notethat FPC cables from CCD packages are not shown in the ﬁgures for clarity.
FEE Installation
The FEE plate is screwed to the base ﬂange (Figure 5.12(e)). Then a FEE board is brought
and 4 (or 6) FPC cables are securely connected to the FEE board (Figure 5.12(f)) . After

5.1. CCD DEWAR

117

(a)

(b)

(c)

(d)

(e)

(f)

(g)

(h)

Figure 5.12: The assembly sequence of the CCD dewar.
that, the FEE board is inserted to the bus connector and the insertion slot which is attachedat the FEE plate and the support bar is ﬁrmly screwed to the FEE plate (Figure 5.12(g)).

118

CHAPTER 5. HSC CAMERA MECHANICS

(i)

(j)

(k)

(l)

(m)

(n)

(o)

(p)

Figure 5.13: The assembly sequence of the CCD dewar.
This process is repeated for all CCDs (26 FEE boards). Finally, all the CCDs are connectedto FEE boards and the FEE boards are rigidly ﬁxed to the FEE plate (Figure 5.12(h)).

5.2. SHUTTER ASSEMBLY

119

Cabling and Vacuum Seal
The whole assembly is tumbled and now CCDs face to the ground (Figure 5.13(i)). The
back assembly is attached to the base ﬂange (Figure 5.13(j)). At this moment, we have aneasy access to the cold straps which are attached to cold heads of the coolers (not shown inthe ﬁgure) and the heat path on the cold plate. The cold straps are screwed to the heat pathﬁrmly to secure the thermal contact. The whole assembly is tumbled again and the cables
to the FEE boards are connected and well-ordered (Figure 5.13(k)). The wall (without
connectors) is covered (Figure 5.13(l)) and cables from FEE boards are drawn from holeson the wall and connected to hermetic connectors (Figure 5.13(m)). The connector plate is
attached to the wall (Figure 5.13(n)) and all the hermetic connectors are attached to thewall (Figure 5.13(o)). Finally, the CCD protection lid is uncovered and the window assemblyis attached (Figure 5.13(p)). The assembly sequence is completed.

5.1.8

Prototyping and Development Schedule

Table 5.9 summaries the development schedule of the CCD dewar.
At this stage, the design of the prototype is completed and the machining factory is working
hard to manufacture the parts. The assembly of the prototype will start from coming April.
After that, we will reﬁne the design based on the result of prototype testing. We will start
the assembly of ﬁnal model from 2010 April and expect that the CCD dewar is ready by
September of 2010.

2009. 02
Design of Prototype complete
2009. 04
Cooler production starts
2009. 04-09 Prototype assemble and test
Vacuum test
Thermal test (using 1 cooler or LN2)
Conﬁrmation of assembly sequence
2009. 09-12 Design of ﬁnal model
2010. 01-03 Production of ﬁnal model
2010. 03
Coolers are delivered
2010. 04-09 Final model assemble and test
Shutter, Filter Exchanger are attached2010. 09
HSC total assembly test starts
Table 5.9: Development schedule of the CCD dewar

5.2

Shutter Assembly

The shutter is one of important components of HSC since it determines the photometricaccuracy of the data. Furthermore, the shutter is the most actively moved component andeasy to be exhaused. The design of the shutter assembly is therefore diﬃcult and must be
done with extreme care.

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5.2.1

CHAPTER 5. HSC CAMERA MECHANICS

Requirements

The shutter is one component of the camera mechanics. The camera mechanics needs to be
detachable and pass through the top frame and the instrument rotator in the case such asscheduled maintenance. Because the shutter has the largest envelope circle diameter among
the camera mechanics components, the package size of the shutter determines the diameterof top frame and can aﬀects the size and the weight of whole unit. From the viewpoint of
reducing the unit weight, the envelope of the shutter should be as small as possible. On theother hand, the shutter must have an aperture larger than 600 mm in diamter so that it does
not block the incoming light. The back focus is 191.9 mm along the direction of optical pathand 60 mm thickness is assigned for the shutter. These space constraints make the designdiﬃcult.
In order to guarantee the uniform exposure time over the focal plane, movement uniformity,
which is deﬁned as the uniformity of the shutter speed in motion, must be less than 1 % p-p.
The variation of the absolute exposure time (i.e. the time between an exposure starts andcompletes) must be accurate to 0.01 sec p-p. Minimum exposure time is targeted to be 1 sec.
All the mechanical or electrical components are designed aiming at 105 cycles for scheduled
maintenance and 106 cycles for operational life. These numbers of cycles correspond to 1 year
and 10 years, respectively. The weight assigned for the shutter is 27 kg including subsystem.
All above requirements, as summarized in Table 5.10, must be taken into account when
we design the shutter assembly.

Packaging constraint
Movement uniformity
Accuracy of exposure
Shortest exposure time
Minimum operational life
Minimum time between scheduled maintenance
Weight

60 mm thickness≥600 mm aperture
1% p-p0.01 sec p-p
1 sec (target)
106 cycles
105 cycles
27 kg

Table 5.10: Requirements for the shutter.

5.2.2

Mechanical Design

There are several types of shutter which are used for astronomical instruments including thesliding door type shutter which is adopted for Suprime-Cam. Among these, the roll type
shutter looks most suitable for our case. The roll type shutter consists of two independentﬁlms which open and close the aperture alternately (i.e. one opens the aperture, and theother comes to close the aperture) and work as a shutter. This type of shutter is generally
made as a compact assembly and we can easily guarantee the unifom exposure time over the
focal plane by adopting this type. Therfore, we adopt the roll type shutter as the shutter
assembly.
Figure 5.14 summarizes the mechanical design of the shutter. Major components arelabeled for reference. The ﬁlm is made of ﬁber glass sheet laminated by black polyvinyl

5.2. SHUTTER ASSEMBLY

121

Film

Ball Screw

SpringDrive Unit

Aperture

RollFigure 5.14: A picture of prototype shutter. Major components are labeled for reference.
Note that the cover plate is taken oﬀ in this ﬁgure.
chloride (PVC). It has 0.3 mm thickness and its transparency is less than 5.9×10−6 atwavelength between 300 and 1,060 nm[1]. One AC servo unit drives two ball screws and one
roller. Blocks on each ball screw are connected to an aluminum bar which holds one edge ofa ﬁlm. The axle of the roller is not directly connected to the AC servo motor because there issome speed diﬀerence between the ball screws and the roller. Instead, we put a spring along
the axle to absorb the speed diﬀerence. This spring also works to apply a pretension of 68 N
to the ﬁlm, which is necessary to wind the ﬁlm properly. AC servo motor is equipped withabsolute encoder so that we do not need extra switches which will be the source of trouble
in the regular operation. An aluminum chassis contains these component inside of 1,030 mm
diameter envelope while having the 600 mm aperture.

5.2.3

Prototyping

Mechanical design of a prototype was started from August 2008. At that moment, the prototype was designed to employ two motors for each ﬁlm; one motor was used to wind/unwind
the ﬁlm and the other was used to drive the bar. The prototype was completed in October
2008 and we started endurance test, tilt test and low temperature test[5].
First, the endurance test for 105 cycles, assuming the interval of the scheduled maintenance,
was carried out. No major trouble was occurred but the spring at the axle had a plastic
deformation after the test. Although the spring met a minimum requirement, we plan torevise the design of the spring so that more reliable operation is realized.

122

CHAPTER 5. HSC CAMERA MECHANICSAt the tilt test, the shutter was tilted at 90 degrees (the bar moves up and down) anddriven 120 cycles continuously. We monitored the motor current to check the axial load andfound that ball screw motors tended to be overloaded when pushing up the ﬁlm. Because
the load is the sum of the weight of ﬁlm and its tension, the load can be controlled byadjusting the tension of the ﬁlm. The load was actually reduced after we adjusted thetension appropriately.
At last, the low temperature test was carried out under the condition of –10◦ C which isassumed as the lowest ambient temperature at the operation. This test was made up of 105
cycles continuous drive and several one shot drives after 1 hour cooling. It was found thatthe diﬀerence of CTE between the ﬁber glass ﬁlm and the metal roller made unexpected
tension to the ﬁlm and resulted in an overload of the ball screw motor. This problem isresolved by (1) assembling the shutter on-site (i.e. under the temperature environment), or(2) making the roller by ﬁber glass.
In conclusion, no signiﬁcant failure was found for the prototype during these tests, however,some parts should be improved to ensure more reliable operation.
A minor problem was that the weight of this prototype exceeded the requirement. So westarted the design to reduce the weight. New prototype adopted one motor for each ﬁlm sothat two motors were taken out and the weight met the requirement. We have just ﬁnishedendurance test of the new prototype and found no major problem so far.

5.2.4

Development Schedule

We have continued the test of new prototype and will complete by March 2009, then start
the implementation for precise control from April 2009. Final model design will be started
from August 2009 and production will be completed in December 2009. We are planningﬁnal test before release and it will be completed in February 2010.

5.3

Filter Exchange System

The ﬁlter exchanger system (FES) provides the possibility to change ﬁlters during the night.
Based on the space available for the system, the FES is separated into two parts: the focal
part attached to the focus unit and the main body attached to the housing of HSC. The
focal part provides the centering mechanism of the ﬁlter while the main body provides the
motions of the ﬁlters.

5.3.1

Requirements

The focal part is severely constrained between the dewar window and the shutter which isonly 55mm along the optical axis. The main body of the FES extends into the light collecting
area of the telescope. It is thus important to minimize the size of the FES especially in thelateral direction. Such requirement limits the possible number of ﬁlters that FES could
accommodate. The current design is able to equip six ﬁlters during the observation. The
space for the ﬁlters can also be used by Shack-Hartmann (SH) optics for the optical alignment
check. To ensure the reproducibility of the SH test, the position reproducibility of the ﬁlters
is set to be smaller than 50 micron and the tilt error should be within 3 arcmin. To minimize
the overhead, the ﬁlter exchange time should be less than 10 minutes with a target time of
5 minutes. Detailed requirements are summarized in Table 5.11.

5.3. FILTER EXCHANGE SYSTEM

Packaging constraint
Number of ﬁlters
Max time between two visits in diﬀerent ﬁlters
Minimum operational life of ﬁlter
Minimum time between maintenancePositioning repeatability
Weight123

55 mm in vertical direction at focal unit
≥ 4 ﬁlters
10 min (goal 5 min)
40,000 cycles
4,000 cycles
50 µm in shift and 3 arcmin in tilt
15 kg (Focal Unit)
300 kg (Main structure including ﬁlters)

Table 5.11: Requirements for FES.

5.3.2

Mechanical Design

Since the FES will be installed separately, the interface and the clear area for the top unit
exchanger are needed. Almost half of the outer region of the FES is used for the mounting.
This limits the possible conﬁgurations of the FES.
The conﬁguration of FES based on the current design is shown in Figure 5.15. Major
components of FES, focal unit, stacker, elevator and tractor, are labeled in the ﬁgure. The
detailed design was carried out taking over the conceptual design conducted at NAOJ[6].

Figure 5.15: Overview of the ﬁlter exchange system. Major components are labeled forreference.

5.3.2.1

The Main Structure of FES

The main body of the FES consists of two identical sets (FES-Opt, FES-IR) installed at
opposite sides of the camera structure. The design of this set is shown in Figure 5.16 (a).

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CHAPTER 5. HSC CAMERA MECHANICS

Each set includes a stacker, a tractor and an elevator. The stacker is for the storage of theﬁlters and is folded close to vertical direction during exposures (Figure 5.16 (c)). Threeﬁlters can be installed in each stacker. For the current design, the eﬀective collecting area
blocked by the stacker is less than 5%. Within the stacker, motorized elevator is design
to select the ﬁlter for use. Four ball screws are used to guide the movement of the ﬁlters
during the selection. Vertical guide blocks are installed to lock the ﬁlters not in use during
operation. The tractor is designed to move the ﬁlter in or out of the optical path (Figure 5.16(b)). Two grippers actuated by DC motors could catch the ﬁlter frame from the near sideedges. The tractor moved along the rails with the power from the geared stepping motorwith precise travel distance. Vertical and lateral guide rollers are equipped to maintain the
direction of the movement. To prevent the jam into the rail gaps, the mass center of the
tractor is designed to be close to the driving rollers. At storage position, the tractors arelocked by a solenoid even when the power is oﬀ. For the ﬁlter assembly, handle blocks arealso designed on the ﬁlter frame for the grippers. The ﬁlter assembly is passively moved by
the tractor with lateral and vertical guider rollers.
5.3.2.2

The Focal Unit of FES

Figure 5.17 shows the schematic view of the focal unit of FES. Four lifting actuator isinstalled on the dewar to lift the ﬁlter to the dewar. When ﬁlter is moved into the focal
unit, the lock pin will lock the ﬁlter at the right position for lifting. The actuator will lift
the lifting frame which holds the ﬁlter assembly to the dewar. The location of the ﬁlter isthen deﬁned by the guide pins and the conic coupling. The conic coupling on the dewar andthe ﬁlter assembly ensures the position and tilt reproducibility. The original idea is to use
the curvic coupling. However, given the size of the ﬁlter, it is diﬃcult to ﬁnd venders that
could produce the curvic coupling of 60cm diameter and thus expensive. The conic couplingis much easier to make and the precision needed can be achieved if enough thickness isprovided. With the manufacturing error of the conic radius around 10 micron, the error oftilt angle will be within 3 arcmin if the thickness of the coupling is more than 11mm for a 12
degree coupling. As mentioned, the vertical distance available for the focal unit is limited.
Although the conic coupling at the dewar side is already occupied and embedded in the wallof the dewar, the maximum thickness for the eﬀective coupling is only 10mm. It is slightly
oﬀ the target value. To overcome the problem, the conic coupling on the dewar side and theﬁlter side will be manufactured simultaneously to ensure better accuracy. A prototype will
be made soon to check the result.
5.3.2.3

Weight Estimation

The focal unit is attached in HSC structure, so the weight constraint is very tight. Lightweight process has been considered for the focal unit. The heavy controller and power supply
has been removed from the focal unit. The current estimation for the weight of the focal
unit of FES is 15.5Kg which is slightly higher than the target value of 15Kg. The steel
conic coupling on the dewar is the heaviest component which is about 5.8Kg. The actuatorsand drivers are 2.46 Kg. The lifting frame is probably the element which weight could be
reduced. On the other hand, the ﬁrst design result shows a total weight estimation of 400
Kg for the main structure. It is much higher than the targeted 300Kg. The three ﬁlters areabout 83Kg in total and the stacker assembly is 69Kg. The current stacker design assumes

5.3. FILTER EXCHANGE SYSTEM

125

(a)

(b)

(c)

(d)

(e)
Figure 5.16: (a) Overview of the main structure of FES. (b) The design of the tractor.
(c) The design of the elevator. (d) The design of the stacker. (e) The design of the ﬁlter
assembly.

the full metal covers of the ﬁlters to prevent the possible dust and other contamination. It is
expected that the weight can be reduced in this part. Also the interface structure between
the HSC and the stacker will be revised to further improve the weight. Lighter materials

126

CHAPTER 5. HSC CAMERA MECHANICS

Figure 5.17: Schematic view of the focal unit.

will be used if the weight is still much higher than the target value after the revisions.

5.3.3

FES Operation Sequence and Operation Time

The operation of FES includes several steps. To insert a selected ﬁlter, ﬁrst, the arm is
unfolded. Then, the desired slot of the stacker is selected for the tractor to grab the ﬁlter
assembly. The tractor pushes the ﬁlter into the focal unit and then leaves the ﬁlter in the
focal unit and retracts. Final step, the focal unit lifts and pushes the ﬁlter to the dewar. To
extract the used ﬁlter, the above sequence is reversed. To exchange ﬁlters, both sequences
are needed. The detailed operation time break is summarized in Table 5.12.
The total time is about 7 minutes. The procedure assumes the ﬁlter lifting happens
simultaneously during the tractor movement. Some of the steps could be combined to shorten
the operation time. For example, the selection and grabbing of the second ﬁlter could bedone during the extraction of the ﬁrst ﬁlter. However, safety checks have to be implemented
from malfunction.

5.3.4

Control and Power System

As mentioned, the whole EFS is actually separated into three pieces. Combined with further constraints on the weight and power/signal connection, the controller system is more
complicated than it should be. To reduce the weight at the focal unit, the master controller
will be located in the control building with Ethernet connection to small slave controllers at
the focus unit and the main structure (Figure 5.18). Each slave controller will control the
motors based on signals from the sensors and limit switches. The master controller is basedon 2.6.25 Linux system to provide the command sockets for the HSC controller. For the
power system, most motors for FES require 24V DC power. The 110V AC power provided

5.3. FILTER EXCHANGE SYSTEM

Extend Arm
Select an unoccupied slot
Tractor moves to grab ﬁlter 1 from parking
Tractor grabs ﬁlter 1
Tractor pulls out ﬁlter 1
Tractor releases ﬁlter 1 and store ﬁlter 1
Tractor moves to parking position
Select ﬁlter 2
Tractor grabs ﬁlter 2
Tractor pushes ﬁlter 2
Tractor releases ﬁlter 2 and store ﬁlter 2
Tractor moves out to parking
Store Tractor
Store Stacker
Retract Arm

127

30 sec20 sec26 sec20 sec65 sec20 sec15 sec20 sec20 sec65 sec20 sec26 sec15 sec20 sec30 secTable 5.12: The detailed operation time break of FES.

from the telescope will be converted by DC regulator on FES. The other option is to place
the convertor in the power box on the top ring. This could solve the heat dissipation issue
and also reduce the weight of the focal unit. The possible voltage drop for this option needs
to be investigated.

5.3.5

Mounting of FES and Loading of Filters (Day-Time Operation)

The FES will be stored at the observation ﬂoor of the dome when rested. A new 1 ton hoist
is suggested to be installed at TUE ﬂoor for the lifting of FES. After lifting to certain point,
FES will be handed over to the unit loader. A new interface will be manufactured to ﬁt the
FES with the unit loader. The unit loader will then install the FES on HSC. However, theunit loader has limited access to the top end. Some interface structure is needed. To install
the FES in one time, the interface need to be powered to adjust the distance of the two setsof FES (Figure 5.19). Two diﬀerent interfaces for FES-Opt and FES-IR will be needed if wewant to have simple structure interfaces. The issue is still under investigation. The loadingof the ﬁlter will also be activated by the controller on FES. As mentioned, each FES set has
a completed power and controller system. The same power and Ethernet connection as on
HSC are provided at observation ﬂoor. It will be used to move the ﬁlter out of the stacker.
An extension rail and support structure will be designed to support the loading/unloadingof ﬁlters. A special cart will be designed to carry the ﬁlters since each ﬁlter assembly is quite
heavy. A protection cover will be included to prevent the possible contamination duringtransportation.

128

CHAPTER 5. HSC CAMERA MECHANICS

Figure 5.18: Connection diagram of the controll system.

5.3.6

Development Schedule

The detailed design of the FES will be ﬁnished by March 2009. Prototypes for three criticalcomponents including the centering mechanism, the stacker and the tractor will be manufactured starting from April 2009. The prototype tests will be ﬁnished in November 2009.
The completed system will be ﬁnished in April 2010.

5.4

Camera Mechanics System

5.4.1

Back Focus Conﬁguration

Figure 5.20 shows the conﬁguration of the back focus. Total back focus 191.6 mm is sharedwith the CCD dewar, the shutter assembly, the ﬁlter exchange system and the lens barrel of
the wide-ﬁeld correcotor. In this ﬁgure, the focal unit of the ﬁlter exchange system is lifted
up and the ﬁlter assembly is pushed to the CCD dewar (i.e. the observing position). The
incident beam, which is plotted in magenta, traces the aperture of the shutter assembly fairly
well. This means that no additional vignetting is occurred by camera mechanics components.
The length between the shutter assembly and the top of the lens barrel is 7.6 mm, which islarge enough to prevent the shutter assembly from touching the lens barrel.

5.4. CAMERA MECHANICS SYSTEM

129

Figure 5.19: The one-time installation sequence of FES.

5.4.2

System Power Consumption

Table 5.13 summaries the maximum and typical power consumption of the camera mechanics
system. We expect to have ∼1400 W heat inside the prime focus unit. This heat must be
taken away by coolant as discussed in the next section.

Component
Cooler
CCD Readout Electronics
Filter Exchanger
Inside PFU
Outside PFU
Shutter
Misc. Electronics

Maximum (W) Typical (W)
900
450∗
1230
380†
217
894
1440
∼500

126‡
334‡
340‡
∼100

Table 5.13: The estimate of maximum and typical power consumption of the camera mechanics. Note: ∗: assuming that coolers work at half power. †: when reading CCDs. ‡:when in stand-by mode.

5.4.3

Heat Exhaust System

To exhaust the system heat, we plan to install the heat exchanger system from Lytron Co.
The selection of the heat exchanger is as follows[4]. (1) First, deﬁne the required heat load
Q, temperature of incoming liquid Tliq , temperature of incoming air Tair and liquid ﬂow rate
F. (2) the initial temperature diﬀerence (ITD) is calculated to be ITD = Tliq - Tair . (3)With the performance capability (Q/ITD) in hand, refer to the thermal performance graphs
provided by the vendor and ﬁnd the appropriate heat exchanger which has a performancecapability larger than the required performance capability.
For our case, the coolant supplied to the prime focus is ethylene glycol. The temperature
is 4 degrees lower than the outside temperature and the ﬂow rate is 5 l/s. Assuming that(1) the outside temperature is 4 degrees (i.e. Tliq =0), (2) we can assign 3.785 l/s (i.e. 1
gallons/s) for the heat exchanger, (3) the required heat load is Q=1400 W, and (4) the130

CHAPTER 5. HSC CAMERA MECHANICS

Figure 5.20: Back focus conﬁguration.
required air temperature is Tair =14◦ C, we ﬁnd that the required performance capability is100. Comparing the thermal performance graphs, Part No. 6320, 6340, 4320, M10-160,
M14-120 and M14-240 can be selected as candidate. Among these candidate, M10-160 is
the most compact and light and we choose M10-160 as the primary candidate. This heatexchanger has a dimension of 426 mm × 269 mm × 55 mm (without fan) and weighs 3.6 kg.
M14-120 is another more powerful but heavier candidate (4.5 kg). Figure 5.21 summarizes
appearances, dimensions and performances of M10-160 and M14-120.

5.4.4

Weight Estimate

Table 5.14 summaries the weight budget and the current estimate weight of the camera
mechanics. The best estimated total weight exceeds the weight budget. We are trying toresolve this issue.

5.4. CAMERA MECHANICS SYSTEM

131

Figure 5.21: Appearances, dimensions and performances of our candidate heat exchanger,
M10-160 and M14-120.

Component
Budget (kg) Estimate (kg)
Dewar
200
210.786
Filter Exchanger
Inside PFU
15
8.82
Outside PFU

(300)

(300)

Shutter
Electronics(Filter)
Total

27
60
(30)
302

27.1
78.134(23.5)
324.84

Table 5.14: Weight budget and current estimate for the camera mechanics

Bibliography[1] Kamata, Y. 2008, HSC Report, HSC20080911 KamataY BlackCurtainTest[2] Komiyama, Y., Miyazaki, S., Kawanomoto, S., Morokuma. T., Nakaya, H., Furusawa,
H., and Tanaka, Y. 2008, SPIE, 7014, 160[3] Komiyama, Y. 2009, HSC Report, HSC20090107 Komiyama CCDemissivity2[4] Lytron Catalog 1997
[5] Uraguchi, F. 2008, HSC Report, HSC20081217 UraguchiF shutterTestReport[6] Uraguchi, F. 2008, HSC Report, HSC20080909 UraguchiF FilterExchgSysCoD

133

Chapter 6
Sensor
6.1

CCD Detector

6.1.1

Requirements for Hyper Suprime-Cam CCD

The requirements of CCD are summarized in Table 6.1. As is adopted in the most of themodern mosaic CCD cameras, we employ the standard format of 4k×2k pixels. Pixel size
can be 13 ∼ 15 µm to realize the relevant sampling of seeing limited image (∼ 0.2 arcsec
for 15 µm pixel). In order to minimize the gap between CCDs in the large mosaic of Hyper
Suprime-Cam, we require to minimize the width of insensitive area smaller than 5 mm along
one short side where the bonding pads are located. The width of other three sides must be
smaller than 0.5 mm.
Small height variation and co-planarity of the detectors’ surface is crucial to keep the goodand uniform PSF over the ﬁeld of view. The system designs of HSC allocates the error
budget of 25 µm Peak-to-Valley for the global height variation. ’Global’ means that theﬂatness between bottom of the package and the light incident surface must be guaranteed
all across the mounted CCDs.
One of the strength of Suprime-Cam is relatively high quantum eﬃciency in longer wavelength, which enabled discovery of many high redshift objects. This was brought by the
deeply depleted CCDs built by MIT Lincoln Laboratory whose depletion layer is about 40µm thick. We want to strengthen the performance even further to save telescope time in thelarge scale survey. This require thicker depletion layer and it could only be achieved by fully
depleted device across the entire silicon thickness.
Photo-generated charges are drifted by electric ﬁeld inside the depletion layer up to the
CCD potential well below the transfering gate. During the transit, the charges suﬀer fromlateral spreading which can degrade the spatial resolution. This charge diﬀusion can beserious for shorter wavelength lights that are absorbed near the surface and thus the charges
travels longer distance. System design allocates the error budget of 0.17 arcsec FWHM forthe image degradation due to the charge diﬀusion. This requires the standard deviation of
the diﬀusion σD < 7.5 µm.
The exposure length is usually constrained by the sky photons particularly in the longer
wavelength. Therefore, higher QE in red makes the possible exposure time shorter. In orderto keep the similar length of the exposure time with Suprime-Cam, full well capacitancemust be enlarged. We require roughly 50 % increase of the full well capacitance comparedwith MIT/LL CCID-20.
135

136

CHAPTER 6. SENSOR

Requirement (-100◦ C)
Format (pixel size)
2048×4096 (15 µm2)
Pixel to Package edge
< 0.5 mm
(Serial register side)< 5.0 mm
Global height variation < 25 µm Peak-to-Valley
QE
400 nm> 45
550 nm> 85
650 nm> 90
770 nm> 85
920 nm> 80
1000 nm> 40
CTE (per pix)
Parallel direction> 0.999995 (1600 e)
Serial direction> 0.999995 (1600 e)
Dark Current< a few e/hour/pix
Charge diﬀusion
σD < 7.5 µm (400 < λ < 1050 nm)
Full well
1 % departure> 150,000 e
Amp. Responsivity> 4 µV/e
Readout noise
150 kHz readout
<5e
Items
Packaging

Measured
0.410±0.025
4.975±0.025
42
87
94
91
78
40
0.999999
0.999998
1.4
7.5
180,000
4.5
4.5

Table 6.1: Speciﬁcations of CCDs and the actual performance measured on the prototype.
All measurements are made at -100 ◦ C.

Speciﬁcation of charge transfer eﬃciency (CTE), dark current and ampliﬁers readout noise
are all equivalent with what were realized on CCID-20 and other science grade CCDs.

6.1.2

Hamamatsu fully depleted CCD

6.1.2.1

NAOJ - HPK collaboration

When we started a conceptual study of Suprime-Cam in 1993, the procurement of large
format CCD was one of the biggest concerns. We at ﬁrst pursued the possibility of building
CCDs in Japan and thus contacted Hamamatsu Photonics which had just started developing
CCDs for science applications for the ﬁrst time in Japan. They provided us with small (∼ 1
cm2 square) thin back illuminated CCDs and we evaluated the devices under the cryogenic
and low light condition. This was the beginning of the collaboration between NAOJ and
Hamamatsu.
In 1997, Hamamatsu had successfully delivered the ﬁrst large format CCD (2048×4096
15 µm pixel three side buttable). The test results showed that the CCD featured excellent
performance in CTE (> 0.999995), dark current (a few pix/hour) and the readout noise(3.5 e for 20 kHz pixel rate) [1]. It was impressive that the defects were quite few (onlyseveral column defects) for such large format devices, which implied the high level of quality
control of the product at Hamamatsu. This front illuminated (FI) CCD was supposed to beprocessed into back illuminated (BI) CCDs to enhance the quantum eﬃciency.
In the meantime, MIT/LL eventually started delivering the BI CCID-20 and the performance was promising. We thus gave up adopting Hamamatsu CCDs to Suprime-Cam

6.1. CCD DETECTOR

137

Figure 6.1: The ﬁrst 2k4k CCD built at Hamamatsu in 1997. This was front illuminated
CCD built with ordinary p type silicon wafer.
CCD Structure
Si Thickness
Vertical clock phase
Horizontal clock phase
Output Ampliﬁers
Package Material

Full Frame Transfer (Can be Frame Transfer with light shield)
200 µm (Can be 100 ∼ 300 µm)
3 phases
2 phases or 4 phases
4 one stage MOSFET on chip and one J-FET on package
Aluminum Nitride
Table 6.2: HPK CCD conﬁguration

because the schedule of Subaru ﬁrst light was too soon. The project with Hamamatsu wasceased for a while and we concentrated on the commissioning of Suprime-Cam
A group at Lawrence Berkeley National Laboratory (LBNL) developed p-channel CCD
on the high resistivity n-type silicon wafer in late 90s. The device is fully depleted overthe entire silicon wafer. Stimulated by the LBNL development and realizing the scientiﬁcimportance of this type of detectors (high QE in red), we resumed the collaboration with
Hamamatsu to pursue the possibility of building the same type of CCDs.
We started the project by building small FI prototypes (1024×256 24µm) and the ﬁnal
outcome was 2k×4k fully depleted BI CCD that saw the ﬁrst light at Subaru in the new
Suprime-Cam dewar in 2008. The progress of the developments was reported in the seriesof SPIE papers [2, 3, 4]. We describe packaging and the performance of the CCD in thefollowing sections.
6.1.2.2

Format

The conﬁguration of the CCD and the block diagram is shown in Table 6.2 and Figure 6.2,
respectively. Image area is clocked by two independent parallel registers so that the bottom
half can be used for storage area if it is masked. Three phase structure is introduced for theparallel register to increase the full well capacitance of the pixel whereas the serial register
remains four phase as is standard in Hamamatsu CCD. The serial register is split into two
section and the ﬂoating diﬀusion output ampliﬁer is implemented on both ends of the section.

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CHAPTER 6. SENSOR

Figure 6.2: Block diagram of HPK FDCCD.
Thus we have 4 ampliﬁers in total. Since the charge can be shifted either direction on theserial register, users can choose the number of output ampliﬁer from 2 to 4. In order to makea space for the routing of two ampliﬁers in the middle, the last 48 lines of parallel area isskewed and the pixels in the area do not have a uniform size. We do not use this area for
imaging.
The wafer adopted for the CCD is made of N-type Si that has very high resistivity (nearly
10 kΩ cm). The bulk CCD wafer is mechanically polished and made thinned to the designed
thickness. It has been demonstrated that the thickness d of 100 µm up to 300 µm is feasible.
Thicker CCD is desired for higher QE but we lose spatial resolution due to the lateral charge
diﬀusion. We decided to employ d of 200 µm which gives allowable level of charge diﬀusion,
σD < 7.5 µm. The CCD has a structure which enables the backside (light incident side)
positively biased to minimize the diﬀusion. See section6.1.3 about the measurement of thecharge diﬀusion.
6.1.2.3

Packaging

CCD package consists of two components: a CCD base and a pin base (Figure 6.3). The
CCD base that mounts CCD silicon chip and has a pin grid array (PGA) at the bottom
side. The CCD base has a narrow slit on the CCD bonding pads so that binding wires are
fed through the slit and wired to the pads on the CCD package. There is inter-connection
between the pads and the PGA inside the CCD base. The CCD base is made of AlN and isbuilt by Kyocera. CCD silicon was ﬁxed by a ordinary die-bond adhesive in the prototype.

6.1. CCD DETECTOR

139

In order to realize better ﬂatness, Hamamatsu switched it to a low viscosity epoxy with Si
shims at the four corners. As a result a ﬂatness of ∼ 10 µ is realized.
The other package component, the pin base made of AlN (Mitsui Mining Material), has
three Ti alignment pins among which two pins have tight tolerance against mounting holeson cold plates for the alignment. This is similar to what we developed for Suprime-Cam[5]. The pin base has a square aperture in the middle for the feed-through of a ﬂex cable.
The uniformity of the thickness of the epoxy resin adhesive is realized by a screen printing
method and the adhesive is cured thermally. Detail drawings of the CCD base and the pinbase is shown in Figure 6.4 and Figure 6.5, respectively.
The ﬂex-cable has, on the end of the CCD side, 3.00 mm staggered arrays of socket which
is customly built by Advanced Interconnections based on their “Peel-A-WAY” scheme using
Type-432 solder preform terminal sockets. On the other end of the ﬂex cable, we adopt
an AirBorn’s 65 pins high density (0.635 mm pitch) connector (NM-225-065-225-TH00).
Smooth insertion is realized by using two side screws. The dimension of the copper tracks
is the following, width = 100 µm, thickness = 18 µm and length = 100 mm. The numberof tracks is 59 in total. We have checked that suﬃcient thermal insulation is realized by theﬂex cable (See Chapter 5).

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CHAPTER 6. SENSOR

Figure 6.3: Structure of CCD package (a) and the appearance (b).

6.1. CCD DETECTOR

141

Figure 6.4: HPK CCD and the AlN CCD base. Drawings with cyan shows the dimensionof CCD silicon. The ceramic base is made larger by 50 (± 25)µm from the silicon. Thus,
the insensitive region on the package is 4.975 mm wide along the horizontal register side and0.410 mm along the other three sides.

Figure 6.5: Detail drawing of AlN Pin base

142
CHAPTER 6. SENSOR

6.1. CCD DETECTOR

6.1.3

143

Performance veriﬁcation

In this section, we showed results of performance tests done at NAOJ detector laboratory.

6.1.3.1

Quantum Eﬃciency

We had expected QE in longer wavelength thanks to the thick depletion layer. Yellow line in
Figure 6.6 represents the QE data of Hamamatsu CCD which we adopted for Suprime-Camin 2008 whereas green line shows the QE of MIT/LL CCID-20 that had been used for nearly
10 years in Suprime-Cam. It is demonstrated that QE in z’-band is improved by about 40%.
At the same time, we noted that the QE in B-band is 10 % lower compared with CCID-20’s
QE. Hamamatsu worked on the improvement of the backside (light incident side) charging
process and eventually worked out better way. Red line in the ﬁgure is the result of the
improvement which we expect on CCDs built for Hyper Suprime-Cam. This indicates that
improved Hamamatsu CCD has higher QE than CCID-20 in all the ﬁlter band except uband. Open circles show the speciﬁcations for HSC CCD at each wavelength and we conﬁrm
that all the specs are mostly met.
On the other hand, blue line represents QE of E2V (CCD42-90) adopted for CFHT MegaCam. Hamamatsu CCD has three times higher QE in z’ band but the QE in B is 30 % lower
and QE in u is even lower. This clearly shows the diﬀerence of the characteristics of eachCCD.

6.1.3.2

Charge transfer eﬃciency

Charge transfer eﬃciency is measured using X-ray. Figure 6.7 shows location of pixel (=number of transfer) versus pixel value of the X-ray events. Since the CCD is thick, the charge
is diﬀused to neighboring pixels. We collect these signals from at most three neighboring
pixels that have signals above a certain threshold.
Two visible bands represents the emission lines of the 55 Fe source (5.9 keV and 6.2 keV).
Slope of the bands (α) is a estimator of charge transfer eﬃciency. The CTE η is calculated1
as η = (N α) N where N is the number of transfer. The estimated parallel and serial CTE is
0.999999 and 0.999998, respectively, which is better than the HSC speciﬁcation (> 0.999995).

6.1.3.3

Dark current

Figure 6.8 shows the temperature dependence of CCD dark current. At -100◦ C it is only<∼ 3 e/hour/pix that qualify the HSC speciﬁcation. It is known that QE in the longer
wavelength increases when the temperature gets higher because the number of phonon toassist the charge to overcome the band gap increases.
According to our measurements QE at 1 µm increases from 40 % to 45 % when we change
the temperature from -100 to -80 ◦ C where dark current of 40 e/hour/pix is expected. Thiscan be negligible in the broad band imaging especially in red ﬁlter because the sky brightness
is very high and variable. Therefore, it will be interesting to prepare a high temperaturemode to pursue the ultimate depth in the redder band.

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CHAPTER 6. SENSOR

Figure 6.6: Quantum Eﬃciency (QE) comparison of several brands of CCDs. All the QE
were measured at NAOJ detector laboratory except E2V, which was measured at CFHT.
Transmission curve of Subaru ﬁlter set (u and z’) are also superimposed.

6.1. CCD DETECTOR

Figure 6.7: Result of X-ray test to measure charge transfer eﬃciency

145

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CHAPTER 6. SENSOR

Figure 6.8: Temperature dependence of dark current
6.1.3.4

Best Focus Position

The absorption length of 1 µm photon is about 320 µm (-100 ◦ C) which is comparable withthe thickness of the CCD. Therefore, the best focus position, where the minimum size ofthe image is realized, is deep inside the silicon. Kawanotomo [6] carried out a simple image
simulation. Photons of F/2 cone are generated and incident on Silicon. The position of theabsorptions is determined based on the absorption length (Figure 6.9). By changing theheight, best focus position is searched. Table 6.3 summarized the result in the case of Si
thickness of 200 µm. At the wavelength 700 nm or below, absorption occurs right below thesurface whereas at λ = 1000 nm the best focus position is nearly half way of the thickness.
Table 6.3 also shows the RMS spread of the image, which is very small compared withthe 15 µm pixel size. This is because silicon has high refractive index, the beam becomes
suﬃciently slow (> F/7) inside the silicon. Actual image size is determined by the charge
diﬀusion that is argued in the following section.
6.1.3.5

Eﬀect of Charge Diﬀusion

The expected charge diﬀusion is σD = 7 µm for the transverse distance of 200 µm at-100◦ C [7]. In the simulation, we shift the charge position based on this statistics andrecord it. The diﬀusion is scaled by the transverse travel distance, dt , inside the silicon as
7 × dt /200 µm [6].
The result is summarized in Table 6.4. At shorter wavelength the image size is determinedmostly by the charge diﬀusion. As the wavelength becomes longer, image size gets smaller6.1. CCD DETECTOR

147

0
-20
-40
-60
-80
-100

40
20
-40

0
-20

0

-20
20

40

-40

Figure 6.9: Position of the absorption of F/2 beam photons (λ =900 nm) inside silicon. The
thickness of silicon is set at 100 µm. Focus position is set at 35 µm below the surface whereminimum image size is realized.

λ [nm]
700
800
900
1000

focus pos. [µm] σ [µm] FWHM [”]∗
5.9
0.20
0.006
14.3
0.48
0.015
43.0
1.36
0.04
89.7
1.98
0.06

Table 6.3: Best focus position from the incident surface in each wavelength. The Si thicknessof 200 µm is adopted here. ∗ Pixel scale of Suprime’s is considered.

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CHAPTER 6. SENSOR

λ [nm]
700
800
900
1000

focus pos. [µm] σ [µm] FWHM [”]∗
5.6
6.9
0.21
19.2
6.6
0.20
43.8
5.8
0.18
84.7
4.5
0.14

Table 6.4: Best focus position and the image size in the simple image simulation. Charge
diﬀusion is considered (See text). ∗ Pixel scale of Suprime’s is considered.

because the transverse travel distance tends to be shorter. Note that position of the best
focus is almost the same as the results shown in Table 6.3.
Kamata et al. (2008) measured the image blur due to the charge diﬀusion [4]. Pinhole of 10 µm diameter is reduced by 1/10 and is projected on the CCD using objectivelens(Figure 6.10). The lens NA of ∼ 0.25 is selected to mimic F/2 optics of the Subaru
prime focus. The size of the image is measured by simple gaussian ﬁtting implemented on
IRAF/imexam. Note that the imaging system is fairly undersampled considering the pixel
size of 15 µm. We checked the imexam ﬁtting results of generated gaussian images withvarious sizes. It turned out that the error of imexam result is within 5 ∼ 10 % comparedwith the original σ down to the one third of the pixel size as long as the image center isclose to the center of the pixel. Image and pixel alignment can be easily made before themeasurement by balancing the values of neighboring pixels.
When we set the backside bias voltage of > 40 V, the PSF is smaller than half of the
pixel in both blue and red wavelength (Figure6.11). The blue and red dotted line showsthe expected image size that are calculated from the image simulation assuming the charge
diﬀusion of σD = 7 µm over 200 µm thickness. The measurement is consistent with what
was expected from the simulation.
The image size was also measured at the telescope during the commissioning run of new
CCDs (Aug 1, 2008 HST). Figure 6.12 shows position dependence of the object size (FWHM
in arcsec) in z’ band of 10 second exposure. The image size is well smaller than 0.4 arcsec
which is the best seeing that we have ever seen on Suprime-Cam. This plot clearly demonstrates that the image smear due to the charge diﬀusion in this wavelength is acceptable.
One also note the gradual increase of the image size toward the larger pixel position acrossthe multiple CCDs. This might be due to the miss-alignment of the optics that required
further investigation.
The smallest image size in B-band that we noted so far is 0.48 FWHM [arcsec] (Aug 24,
2008 HST). Since it is relatively rare to have better natural seeing in shorter wavelength, theperformance veriﬁcation has not yet been done below this seeing. However, as is shown in
Figure 6.11, the wavelength dependence of the image blur should not be signiﬁcantly large.
6.1.3.6

Linearity

CCD system linearity was measured and the result is shown in Figure 6.13. Top panel shows
the output response curve (average of 10×10 pixels) where the light yield is controlled by
the exposure time. Bottom panel represents the residual from the linear best ﬁt obtained in6.1. CCD DETECTOR

149

Figure 6.10: Experimental setup to measure the charge diﬀusion
the middle of the data (around 105 e). The results suggests that the system is linear withina percent level up to 180,000 electrons. This is well beyond the HSC speciﬁcation.
6.1.3.7

Readout noise

Figure 6.14 shows the result of readout noise as a function of pixel rate. The rate wasdetermined by the integration time of correlated double sampling circuit and the serial
clocking pattern are common to all the measurements. The noise is measured by RMS
of pixel values on the horizontal overscan region. Suﬃciently low noise is realized for ourapplication (∼ 5 e at 150 kHz readout).
6.1.3.8

Residual image

During the course of the R&D phase of the CCD, we noticed the existence of residual image
after extremely strong light are illuminated. Figure 6.15 shows the example of the residual
image seen after the saturated star imaging.
This can be explained by delayed charge (hole in our case) emission from the silicon surface
traps (CCD gate side) and the countermeasure has been established as the following. Everytime we readout the image, we make the back bias voltage nearly zero and make all theparallel gate to positive (a bit high more positive compared with normal clocking voltage)and keep this voltage one second long. This way, the charges in the surface trap moves tothe n+ channel stop between pixels (this is usuall called “pinning”) and the trapped charges

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CHAPTER 6. SENSOR

Figure 6.11: Charge diﬀusion measurement

Figure 6.12: Object size as a function of ﬁeld position x in pixel. z’ band 10 second exposure
at the elevation of 74 deg. There are 5 chunks of data points and each chunk shows theresult of two overlapped CCDs. along y axis.

6.1. CCD DETECTOR

Figure 6.13: Linearity test result using MFront2 readout electronics.

151

152

CHAPTER 6. SENSOR

Figure 6.14: Readout noise as function of pixel rate.

are erased. We call this clock sequence “erase”. If erase is applied, we conﬁrmed that noresidual like the one in Figure 6.15 image is not observed.

6.1.4

Procurement and test plan

HSC requires 118 CCDs at least and we try to procure 125 CCDs in total including thespares. The procurement plan is shown in Table 6.5. The delivery pace is roughly 10 CCDs
per month and Hamamatsu agreed with the pace. We have to improve the eﬃciency of our
test bench at our lab.

1 2
2009
2010 10 13

3 4 5
5 10 10
12 0 0

6
10

7 8
10 5

9 10 11 12
10 10 10 10

Table 6.5: Hamamatsu CCD procurement plan6.1. CCD DETECTOR

153

Figure 6.15: Example of residual image of bright stars. Two successive images are shown inthe middle and on the right that are taken after the image with saturated stars (left).

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CHAPTER 6. SENSOR

Figure 6.16: Overhead time to read an image

6.2

Front-end Electronics (FEE) and Back-end Electronics (BEE)

6.2.1

FEE and BEE

HSC is planed to use 116 pieces of 2k × 4k CCDs. Although the number of CCDs installed
to HSC is ﬁxed, the requirement of the readout electronics is considered to handle 128 CCDs
for the maximum case [13].

6.2.1.1

Requirements

The performance of the fully-depleted CCD manufactured by Hamamatsu Photonics K.K.
is summarized in the previous section. The values described in the section are obtained withthe readout speed of 130 kpixel/sec. This readout performance is also our target for HSC,
but slightly worse performance can be acceptable. The requirements for the HSC readout
electronics is listed in Table 6.6.
HSC will have two readout modes. One is the standard readout mode, and the other is theslow readout mode which achieves the same readout performance described in the previous
section.
The readout time of the standard mode is 10 seconds and that of the slow mode is 20
seconds. The total overhead times including erase, wipe, and shutter driving time to readan image become 19 seconds and 29 seconds, respectively (Figure 6.16).
The size of the image data is huge compared to that obtained by the traditional astronomical instruments. It is 2.31 GByte including blank and over-scan regions, which is morethan 10 times larger than that of Suprime-Cam. A frame memory where the whole image
data is stored temporarily is needed to guarantee the real-time control of the CCD, and acommunication capability of the data rate 237 MByte/sec is required to sustain 10 seconds
readout continuously.

6.2. FRONT-END ELECTRONICS (FEE) AND BACK-END ELECTRONICS (BEE) 155

Number of CCD
Image Format

1284272 x 2272
(including 8 blank, 48 over-scan row & line)
Total Signal Output
512 ports
Analog-to-Digital Conversion 16bit
Data Size
2.31 GByte/frame (2427MByte/frame)
Readout Time
Standard Mode
10 seconds/frame
Slow Mode
20 seconds/frame
Pixel Rate (calculated from the readout time)Standard Mode
243 kpixel/sec
Slow Mode
122 kpixel/sec
Data Transfer Speed
237 MByte/sec at the standard readout mode
Values after installing onto the prime focus
Readout Noise≤10 e by the slow readout mode
Full Well Capacity≥150,000 e by the slow readout mode
Linearity≤±1 %Table 6.6: Requirement for the HSC readout electronics

6.2.1.2

Overview

The CCD readout electronics for HSC can be divided into two parts; one is the front-end
electronics (FEE) placed in the CCD dewar, and the other part is the back-end electronics(BEE) placed outside of the dewar in the prime focus unit (Figure 6.17).
Since FEE is installed in the vacuumed dewar, the heat generated on the FEE board can
not be directly spread away into the air. The FEE board is built on an aluminum core PCB.
The heat generated in electronics parts is transferred through the aluminum core, and somestructural parts of the dewar which support the FEE board. Then, it is radiated from thedewar walls into the air in the prime focus unit. Please refer to the section 5 for more detail.
The FEE contains all of analog circuits to readout CCDs including clock drivers, preamps,
correlated-double-sampling (CDS) circuits and ADCs. BEE consists of digital circuits including a clock sequencer, a frame memory, and an Ethernet block to communicate with aremote host. The FEE and BEE are connected only via LVDS. Analog signal lines do not
cross the dewar wall at all. BEE is remotely controlled from the observation room through1 Gbps optical Ethernet. The clock to control CCDs is generated in BEE. The image data
digitized by the ADCs in FEE is temporarily stored in the frame memory of BEE, and thentransferred to the observation room by the 1 Gbps Ethernet.
The 1 Gbps Ethernet module is designed based on SiTCP developed in a ﬁeld of highenergy physics [15]. SiTCP is a technique using a ﬁeld programmable gate array (FPGA)
to achieve real-time data transfer limited by the speciﬁcation of the most commonly usedgigabit Ethernet. A clock sequencer and a frame memory controller are also programmed in
the same FPGA. A DDR2 memory module up to 2 GBytes can be mounted on the board.

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CHAPTER 6. SENSOR

Figure 6.17: Block Diagram of FEE and BEE

The entire CCDs of HSC is controlled by 3 groups of FEE/BEE pairs. Two groups are forscience CCDs and 1 group is for guide CCDs. The number of the boards used in those groups
is diﬀerent, but the same FEE boards and BEE boards are used for the 3 groups. Because
each group has its own optical Ethernet, three pairs of the optical ﬁbers are connected toremote computers located in the observation room.
The functions of FEE are summarized in Table 6.7, and the boards of BEE are listed inTable 6.8.
Two types of the FEE boards are designed along with the arrangement of CCDs on the
focal plane. One is to control 4 CCD on a size of 240 mm × 150 mm board, and the otheris to control 6 CCD on a size of 240 mm × 150 mm board. The board arrangement in the

Bias generator
Clock driver
Preamp
CDS circuit
ADC
Power management

Generating bias voltages by DACs
Driving CCD clocks and biases
AC coupling and preamp
Dual-slope type correlated double sampling circuit
16 bit 500k sampling/second analog to digital converter
Linear voltage regulators, monitoring power voltages/currents, and temperature
Table 6.7: Functions of FEE

6.2. FRONT-END ELECTRONICS (FEE) AND BACK-END ELECTRONICS (BEE) 157

GESiCA Interface (BGI)

Communication board using 1 Gigabit Ethernet withremote hosts. This board also has a clock sequencer
and a frame memory.
Clock Transmission (BCT) Interface board for the front-end board, including serial
communication and sending CCD clock signal
Power (BPW)Managing power supply including power consumption
monitoring and shutdown function
Data Acquisition (BDA)
Interface board for the front-end board. Receiving ADC
data, converting the serial data to parallel, then transfer
to BGI
Table 6.8: Board types of BEE

dewar is illustrated in Figure 6.18.
Three types of circular hermetic connectors: 55 pin, 41 pin and 26 pin connectors are
employed. The required numbers for a half of the science CCDs are 2, 4 and 3, respectively.
The 55 pin connectors are assigned to CCD clocks and communication lines between FEE
and BEE, the 41 pin connectors are assigned to ADC data, and the 26 pin connectors are
assigned to power lines. Figures 6.19 and 6.20 show the connection diagrams of BEE and
FEE, respectively.

6.2.2

Front-end Electronics

The FEE board has all of the analog circuits for CCD control and readout. We have built
and installed a prototype readout electronics to the upgraded Suprime-Cam. The analog
part is called MFront2 [14], and the digital part is called MESSIA5 [11][12]. In this section,
we introduce the readout electronics of the upgraded Suprime-Cam and the performance ofMFront2 with MESSIA5, and then explain the FEE circuit for HSC.
6.2.2.1

Prototyping

The MFront2 for the upgraded Suprime-Cam consists of 10 pieces of 3U euro size cards(160 mm × 100 mm). One is an interface card (MF2 IF), two are power management cards(MF2 PWR), two are bias generating cards (MF2 BIAS), and 5 are signal processing cards(MF2 SIG). One MF2 SIG can handle 8 channels of video signals. And also, 10 preamp
boards (MF2 CCDPRE) are attached next to hermetic connectors outside of the dewar(Figure 6.22), while the former Suprime-Cam has preamp modules in the vacuumed dewar.
Suprime-Cam has enough feed-throughs for CCD signals and a space on the outside of thedewar to attach the MF2 CCDPRE boards. It is not mandatory to place preamp boardsin the vacuumed space. Placing electronics boards outside the dewar makes maintenance of
the instrument rather easier than placing them in a vacuumed space as long as it is possible,
although the signal paths becomes somewhat longer.
MF2 PWR for the Suprime-Cam is conﬁgured to supply +54 V for the back-bias circuit
which can apply the back-bias voltage higher than +50 V to the FDCCDs. The clock

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CHAPTER 6. SENSOR

Figure 6.18: FEE arrangement6.2. FRONT-END ELECTRONICS (FEE) AND BACK-END ELECTRONICS (BEE) 159

Figure 6.19: Connection diagram of FEE

Figure 6.20: Connections diagram of BEE

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CHAPTER 6. SENSOR

Figure 6.21: MFront2 for CCD evaluation

6.2. FRONT-END ELECTRONICS (FEE) AND BACK-END ELECTRONICS (BEE) 161

Figure 6.22: MFront2 for the upgraded Suprime-Cam

and bias voltages including the back-bias generated in MF2 BIAS can be changed from aremote computer. It makes us possible to use diﬀerent bias voltages for the exposure andthe readout.
The MF2 CCDPRE is conﬁgured to amplify CCD signals by a factor of 3.0. MF2 SIG
has 16bit analog-to-digital converters (ADC) with 3.0 V input range and is conﬁgured as
the gain of 1 when the integration time of correlated double sampling circuit (CDS) is set to
2 µs. As the result of this conﬁguration, the input gain of the analog front-end electronics
is 15.1 µV/ADU. The node sensitivity of the FDCCD is around 5 µV/e. With utilizingthese values, the conversion factor is calculated to be ∼3.0 e/ADU. The full-well capacityof the CCD can be measured up to 1.0 V which corresponds to ∼200,000 e with the nodesensitivity. The speciﬁcation is summarized in Figure 6.23.
6.2.2.2

Performance of the prototype FEE

To examine the intrinsic performance of MFront2, it was tested with a dummy CCD signal
which comprises precision 16bit digital to analog converters (DAC) and analog switches. The
linearity was measured with the dummy signal by increasing the signal level. The linearity
error which is the deviation of data from the best-ﬁt line is less than 0.1 % throughout the
input range of the ADC.
The readout noise was measured by connecting the preamp input to GND. Noise generally
depends on the readout speed, especially the integration time of CDS. The measurement
was carried out for various integration times of CDS between 1 µs and 64 µs. The 2 µs

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CHAPTER 6. SENSOR

Figure 6.23: Speciﬁcation of the upgraded Suprime-Cam

integration time corresponds to the gain of 1. When the integration time is 2 µs, the input
noise is about 9 µV with the preamp gain of 4.2. When the integration time is 64 µs, theinput noise is decreased down to 2 µV.
Figure 6.24 shows the readout noise versus readout time in case that the 2k×4k FDCCD isread by Mfront2. Note that the readout time depends on the integration time of CDS, andthe amount of the noise was converted to electron with the conversion factor of 4.7 µ/e. The
dashed line shows the intrinsic noise of MFront2 measured by the above method. The thickline shows the noise value measured with a test chip of FDCCD. This is combined value ofthe readout electronics and the CCD. The thin line represents the intrinsic readout noise of
the CCD calculated by subtracting the intrinsic noise of Mfront2 from the combined noisewith Mfront2 and the CCD (the dashed line).
Figure 6.25 shows the linearity error. The dashed line represents the intrinsic error ofMFront2, and the thick line represents the error combined with the electronics and a test
chip of FDCCD. The linearity error originated from the readout electronics is small enoughcompared to the error of the CCD.
The gain temperature drift of the CDS circuit on the MF2 SIG board is also measured.
MF2 SIG is using the double slope type of CDS circuit which consists of an integration
circuit with a capacitor and a resistor. This type of circuit has a mismatch of the temperaturecoeﬃcient between the capacitor and the resistor. We conﬁrmed that the temperature drift ofthe gain of our CDS circuit is within a few 100 ppm/◦ C without any compensation techniques(Figure 6.26). It corresponds to a few ADU per 10 ◦ C which is comparable to the readout
noise of Suprime-Cam. Since the temperature change atop of Mauna Kea is expected to be6.2. FRONT-END ELECTRONICS (FEE) AND BACK-END ELECTRONICS (BEE) 163

Figure 6.24: Noise measurement for MFront2

Figure 6.25: Linearity error measured for MFront2

164

CHAPTER 6. SENSOR

Figure 6.26: Gain temperature drift measured for MFront2

less than 10 ◦ C during a night, the eﬀect of the gain temperature drift can be negligible.
Gain temperature drift less than 100 ppm/◦ C was achieved with an active compensation
circuit to prevent the mismatch of the temperature coeﬃcients. However, it seems that suchan active compensation circuit is not needed for Suprime-Cam and HSC.
6.2.2.3

FEE circuit for HSC

Each component of FEE is designed based on Mfront2. We introduce the clock driver, thesignal processing circuit, and the board management for HSC hereafter.
Clock driver circuit
The clock driver circuit is shown in Figure 6.27, and the CCD signal processing circuit isshown in Figure 6.29. Our analog circuit is based on the SDSS photometric camera [8].
Both of the clock driver and signal processing circuits are well suited for a large mosaic
CCD camera, because current consumption can be kept low in addition to good readout
performance.
Figure 6.27 shows the clock driver circuit. One channel of the clock driver circuit consistsof 2 op-amps, 2 large capacitors, and 1 analog switch. The op-amps drive static voltages ofhigh and low levels of the CCD clock each. This means a high speed driver op-amp to drive
large capacitive CCD electrodes, which consume current, is not needed.
Bias voltages are generated by a 16 ch multiple output DAC with bipolar outputs. Those
outputs of the DAC are converted to the required voltages for the CCD. Serial clocks, output
gate, and back bias gate which do not need current are connected to each CCD in parallel.
Parallel clocks and drain biases which need current are buﬀered for each CCD in contrast toserial clocks (Figure 6.28).

6.2. FRONT-END ELECTRONICS (FEE) AND BACK-END ELECTRONICS (BEE) 165

Figure 6.27: Clock Driver Circuit

Figure 6.28: Bias Current Flow

166

CHAPTER 6. SENSOR

Figure 6.29: Preamp and CDS Circuit

Figure 6.30: ADC: AD7686 manufactured by Analog Devices Inc

Signal processing circuit
Figure 6.29 shows the signal processing circuit. A CCD output signal is fed to an AC coupling
circuit with DC restoration to subtract the oﬀset voltage before preamp. The signal is thenpre-ampliﬁed by a factor of 3, and processed by a dual-slope type of CDS circuit. Finally,
the analog signal is digitized by a 16 bit ADC.
The ADC we chose is AD7686 manufactured by Analog Devices Inc (Figure 6.30). It isa very high dynamic range (92.5 dB @ 20 kHz), small size (3 mm × 3 mm) and low power
consumption (3.75 mW @ 100 kSPS) CMOS analog-to-digital converter. The throughput
is 500 kSPS. Furthermore, the output format of the ADC is in serial. Four outputs of the
ADC can be daisy-chained. Therefore, one CCD with 4 outputs has one serial output. In
this case, the throughput becomes 344 kSPS with a 50 MHz serial data clock, which can becalculated by the following equation.
FS = 1/(TCOM V + (TSCK × 16 × NADC + 1))(6.1)

FS is the sampling frequency, TCOM V is the conversion time, TSCK is the serial data clock
period, and NADC is the number of ADC.
Finally, the serial data is transferred to BEE via LVDS, and then stored temporarily inthe frame memory of BEE.

6.2. FRONT-END ELECTRONICS (FEE) AND BACK-END ELECTRONICS (BEE) 167

Figure 6.31: Power Consumption with the saturated signal in the signal processing part

Board and power management
Low dropout linear voltage regulators are equipped on each FEE to supply the power toanalog circuits. The reasons to use the voltage regulators are to protect the FE board from
unexpected miss-connection of the cables, to compensate the voltage drop between the power
supply and the FE board, and to equip the FEE board with shutdown function. Because
about 10 % additional power dissipation is expected on the regulator, these regulators shouldbe placed near the board edge to exhaust the heat easily.
All of the power voltages, current, and several points of the board temperature are monitored by a CPU periodically. Once unexpected values are observed, it shuts oﬀ all of theregulators.
The CPU is primarily needed to communicate with BEE. Asynchronous communication
method is used for the purpose. The CPU controls DACs to generate bias voltages and
ADCs to monitor power lines according to incoming commands from BEE.
The CPU can go into the sleep state during the CCD readout to avoid interference toanalog circuits. However, we expect the noise from the CPU to analog circuits is negligible
based on our experience of the prototyping described above.
6.2.2.4

Thermal Design

We place the FEE board in a vacuumed environment. The heat from the electronics can
not be taken away by the ambient air. We use aluminum core circuit boards to transfer theheat eﬃciently from the board.
The heat generated on the FE board mostly come from the signal processing part. We
estimated the power consumption of the signal processing circuit for 4 channels (1 CCD)
during stand-by at 0.35 W/CCD, and during readout time with saturated signal at 0.61
W/CCD. One channel of the signal processing part has three op-amps. The op-amp for the
preamp and the one for the integration-amp consume 35 mW each in the maximum case
(Figure 6.31).
The temperature diﬀerence between the package of an op-amp and the FE board edge
where the heat is exhausted to the dewar wall is divided into three parts. They are thediﬀerence between the package and the copper layer, the diﬀerence between the copper and

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CHAPTER 6. SENSOR

Figure 6.32: FEE board Thermal Design 1

the aluminum core of the FE board, and the deference within the aluminum core (Figure6.32).
The temperature diﬀerence between the package and the copper layer depends on the
cross section of the pins of the packages. Because at least 2 pins are connected with goodthermal conductance to the power planes, the temperature diﬀerence in the worst case canbe estimated with the cross section of the 2 power pins. Because a dual type op-amp in a
SOP package is used, the package could become a 70 mW heat source. The temperaturerising on the SOP package from the copper layer is calculated to be 2.5 ◦ C (Figure 6.33).
The temperature diﬀerence between the copper layer and the aluminum core was estimatedas follows. The thickness of the epoxy between the copper layer and the aluminum core isassumed to be 0.5 mm. The amount of the heat generated on the FE board for a CCD,
which includes the heat from the signal processing circuit, the clock driver circuit, and theboard management circuit, is assumed to be 1W/CCD. The area of the circuit for a CCD is
assumed to be 60 mm × 150 mm. If the temperature on the copper layer as well as aluminum
core is distributed uniformly, the temperature diﬀerence between the copper layer and thealuminum core can be estimated to be 0.18 ◦ C with the thermal conductivity 0.3 W/(mK)of glass-epoxy.
Finally, the temperature distribution on the aluminum core can be estimated by adding

6.2. FRONT-END ELECTRONICS (FEE) AND BACK-END ELECTRONICS (BEE) 169

Figure 6.33: FEE board Thermal Design 2

the 1 W/CCD heat inﬂow from a side. The temperature diﬀerence between a side and theother side with a ﬁxed temperature can be estimated to be 5.3 ◦ C with the aluminum coreof 2 mm thickness (Figure 6.34).
The total temperature diﬀerence between a package and the edge of the FE board isestimated to be 8.0 ◦ C in the maximum case. According to the current design of the dewar(Section 5), the temperature of the package is calculated to be slightly lower than 50 ◦ C.
The operating temperature of an op-amp is speciﬁed in the absolute maximum rating section of the data sheet. The junction temperature of the op-amp of the signal processing partis speciﬁed as below 150 ◦ C, and the operating temperature of the same op-amp is speciﬁed
as below 125 ◦ C. The operating temperature of most commercial op-amps is typically speciﬁed under 85 ◦ C. We conclude that the FEE board is operated under the safe temperature
condition.
6.2.2.5

Consideration of the guide CCD electronics

The same Hamamatsu CCD is used for the guide CCDs, but the frame rate is faster thanother science CCD. Only 200 × 200 pixel window is needed for the purpose. The same pairof FEE and BEE is used for the guide CCDs. But the gain of the preamp and CDS might
be changed. The readout time of the guide CCDs consists of the time to wipe the CCD andreadout the CCD. The time to readout the 200 × 200 window is negligible compared to thetime to wipe the whole CCD. The time to wipe the CCD is limited by the driving capability
for the parallel clocks of the CCD. The 3 phase electrodes of the CCD parallel clock haveabout 150 nF of capacitance, and the voltage swings 10 V peak-to-peak. Therefore, the 3

170

CHAPTER 6. SENSOR

Figure 6.34: FEE board Thermal Design 3

Figure 6.35: Guide CCD readout speed

seconds of the wipe time results in 66 mW heat dissipation on the driver package. Thus thewipe time around 3 seconds becomes the limitation to readout the guide CCDs.

6.2.3

Back-end Electronics

6.2.3.1

Overview

The back-end electronics (BEE) provide the interface between the front-end electronics(FEE) and the DAQ system.
Figure 6.36 shows a schematic of the readout system for the HSC. The FEE and BEE are
integrated into the camera, while the DAQ system is located remotely from the camera. BEE
and DAQ communicate over Gigabit Ethernet using Transmission Control Protocol (TCP)and User Datagram Protocol (UDP) as parts of the TCP/IP protocol suite. Ethernet andTCP have been widely used and implemented in a variety of commodity products, and
TCP/IP is the de facto standard network protocol and is implemented in many operating
systems (OS). The DAQ system also employs TCP/IP and Ethernet as a common interface.
Thus, adopting TCP/IP and Ethernet provides high connectivity between BEE and theDAQ system. The main tasks of BEE are CCD control, storage of image data, and transferof image data to the DAQ system within the required readout time. The DAQ system
controls FEE by sending commands to BEE for image data readout. The commands areprocessed by a TCP/IP processor. A CCD readout controller generates readout clocks for
the CCDs. The output image data of the CCDs are converted to digital signals by ADCs in6.2. FRONT-END ELECTRONICS (FEE) AND BACK-END ELECTRONICS (BEE) 171

Figure 6.36: Schematic of the HSC readout system

FEE. The converted image data are transferred to BEE and stored in a frame memory. The
stored image data are read from the memory and transferred to the TCP/IP processor by
a command issued from the DAQ system, and the image data are ﬁnally transferred to the
DAQ system [16][10].
6.2.3.2

Requirement

The requirements for BEE are as follows:• Small size and light weight• Low power consumption• Fast speed data transmission• Low readout noise
BEE is integrated into the prime focus unit. To reduce the load of the camera, the size
should be as small as possible and the weight should be as light as possible. To avoid theinﬂuence to seeing, the power consumption should be small. This also contributes to reduce
size and weight. BEE must be capable of processing one 2.31 GBytes frame of data andtransferring it within 10 seconds. The required processing rate is 1.85 Gbit/s. The CCDs
are controlled with readout clocks generated by BEE for image data readout. Jitter of thereadout clock makes readout noise of image signals [12]. The jitter should be smaller than
Poisson noise of the CCDs at near the zero signal level.
Issues with the development of this BEE are reduction of weight, size, and power consumption while maintaining high-speed data transfer capability. Generally, TCP/IP is processedby software making use of standard OS. A powerful hardware and CPU are required to process these protocols at gigabit rates, and thus power consumption is of the order of several
tens of watts and the size and weight are both large (e.g., a single/double height 6U-Eurocard). To overcome these diﬃculties, we have adopted a hardware-based TCP/IP processor
called SiTCP. As the logic size of the processor is small enough for implementation on asingle Field Programmable Gate Array (FPGA), we can implement circuits for all the main

172

CHAPTER 6. SENSOR

Figure 6.37: Block diagram of the planned BEE system unit

tasks BEE on a single FPGA. This allowed us to design a small and lightweight system thathas high-speed data transfer capability.
Figure 6.37 shows a block diagram of the planned BEE system unit using the 3U Euro-card
system. To satisfy the requirements, three system sets of this are used. Two units are used forprocessing of science data, the signals of a half of the science CCDs are processed by one unit.
Other unit is used for guiding. The system unit consists of seven modules: a Readout Module(GESiCA), two Back-end Clock Transmission boards (BCT), two Back-end Data Acquisitionboards (BDA), and two Back-end Power boards (BPW). GESiCA is the main module in thissystem and is mounted on a Back-end GESiCA Interface board (BGI). GESiCA has thefollowing functionalities: this module generates readout clocks for the CCDs, stores image
data into the Frame Memory, and transfers these data to the DAQ system via Gigabit
Ethernet. The BCT converts signaling standards of digital signals: from/to LVDS to/from
CMOS 3.3 V. The BDA controls the ADCs in FEE with a Serial Peripheral Interface (SPI),
and then transfers converted image data to GESiCA through LVDS signal lines.
6.2.3.3

Prototyping

Aim
To study feasibility of BEE for HSC, we have developed a prototype of GESiCA and measured
its performance.
Figure 6.38 shows a photograph of the prototype of GESiCA seen from the front. The
module has the following features: small size, light weight, low power consumption, gigabit
data transfer rate, and a large memory capacity. The module measures 149 mm × 79
mm × 12 mm and weighs 100 g, compatible with the Common Mezzanine Card (CMC) size
speciﬁed as IEEE P1386. Generally, CMC are designed to be plugged into a slot on the host’s
motherboard. BEE employs a host board of 3U Euro-card size. GESiCA is independent of
the system of host modules. Hence, various systems can be used in diﬀerent phases ofdevelopment. For example, to evaluate the function of CCD readout, the current SuprimeCam system can be used. This is advantageous for development. The power consumption of
the module is less than 15 W. One Gigabit Ethernet port of 1000BASE-T, with Unshielded
Twisted Pair cabling (UTP), is used for image data transfer to the DAQ system. The
ﬁnal version of the module is expected to use 1000BASE-SX. However, at present, we have

6.2. FRONT-END ELECTRONICS (FEE) AND BACK-END ELECTRONICS (BEE) 173

Figure 6.38: Front view of the prototype of GESiCA

adopted 1000BASE-T for ease of debugging the module and the system. One DDR2-SDRAM
Small Outline Dual In-line Memory Module (SO-DIMM) with capacity of up to 2 GBytes is
employed as the frame memory.
Figure 6.39 shows a photograph of the prototype of GESiCA seen from the top. The
module consists of three main devices: a Gigabit Ethernet Physical Layer Device (PHY)(DP83865DVH; National Semiconductor Corp.), an FPGA (XC5VLX50-FFG676-2I; Xilinx
Inc.), and an SO-DIMM. This design includes no CPU. 1000BASE-T Gigabit Ethernet
is used for communication between BEE and a computer in the DAQ, transferring theimage data and exchanging commands and status. The FPGA processes network protocols,
generates CCD clocks, and controls the SO-DIMM as the frame memory. Signal lines between
BEE and FEE are connected via the Pn4 in CMC connectors. All pins of the connector areconnected directly to the FPGA.
Figure 6.40 shows a block diagram of the prototype. The PHY device processes an Ethernet
layer protocol, clock recovery, data encoding, and decoding. The PHY and the SiTCP
are connected with the standard interface, Media Independent Interface (MII) speciﬁed byIEEE802.3. The SiTCP processes TCP packets for image data transfer and UDP packets forslow control. Internet Control Message Protocol (ICMP) and Address Resolution Protocol(ARP) packets for network management are also processed by the SiTCP. The Readout
Controller generates CCD readout clocks with a programmable sequencer and clock patterns
stored in a clock pattern memory. A program for the sequencer and clock pattern data are
set by a PC of the DAQ system with the slow control mechanism of the SiTCP. After the
Readout Controller is started, image data are acquired from ADCs. The image data are
stored in the frame memory by the frame memory controller. Generally, DDR2-SDRAM
devices run at a clock frequency of 125 MHz or higher. We use the frame memory at aclock frequency of 50 MHz to reduce power consumption. Thus, the maximum bandwidth

174

CHAPTER 6. SENSOR

Figure 6.39: Top view of the prototype of GESiCA

between the frame memory and the FPGA is 6.4Gbps. The frame memory controller reads
back image data from the memory and transfers these data to SiTCP when a TCP link isestablished. The image data received by the SiTCP are transferred as TCP packets to the
DAQ system.
Tests and results
We measured throughput and estimate readout noise generated by the prototype of GESiCA.
Throughput
Setup:
Figure 6.41 shows a block diagram of a measurement setup of the throughput, which consistsof three parts: GESiCA, a Test Data Generator, and a DAQ PC. We construct this systemusing two GESiCA boards mounted on a test board supplying DC power, +5V and +3.3
V. One of GESiCAs is a device under the test. Another is used for an emulator of an
ADC in FEE, which generates test data. Its circuit of the generator is integrated into an
FPGA. The PC has Gigabit Ethernet ports, which is a Dell Poweredge 2900 Quad Core
Xeon E5410 (2.33 GHz, 2×6 MB L2 Cache, 1333 MHz FSB) running CentOS 5 64 bit. In
this test, we used this DAQ PC is connected to GESiCA directly with a crossover cable.
GESiCA and the DAQ PC act as a TCP server and a client respectively. To start the test,
a user should let the DAQ PC send a command to GESiCA. After receiving the command,
the readout controller in GESiCA sends a test start command to the test data generator
through the CMC connectors. Once the test data generator receives the command, it sends
16-bit count-up test data and 2 bit control signals to GESiCA via the CMC connectors.
The signal standard of these signals is CMOS 3.3 V. The test data generator stops after6.2. FRONT-END ELECTRONICS (FEE) AND BACK-END ELECTRONICS (BEE) 175

Figure 6.40: Block diagram of the prototype of GESiCA

sending 2 GBytes test data. The test data comes into the frame memory controller andthen is temporarily stored in a DDR2 SDRAM SO-DIMM. The test data are read from the
SDRAM and sent to the DAQ PC through Gigabit Ethernet 1000BASE-T. A generating
rate of the test data can be set to various values.
Result:Fluctuation of the transfer throughput was measured at the required value, 0.8 Gbps. We
have developed software on the DAQ PC for this test using socket programming, which
records the time every time when 16.77 MBytes of data from GESiCA are transferred.
Generally, a special conﬁguration for the socket is required in development for high-speed
readout. There is no special conﬁguration in our program. The throughput can be calculated
from the recorded time-stamps. Figure 6.42 shows the measured result. The X-axis is theelapsed time and Y-axis the measured throughput. The measured throughput from GESiCA
to the DAQ PC is 0.8 Gbps, which is same as the input rate. The result satisﬁes therequirement for the throughput and shows that the throughput is stable throughout thereadout process.
The throughputs at various input rates were also measured. We measured a transfer
period which GESiCA spent to send test data of 2 GBytes size to the DAQ PC. From thetransfer period we then calculated the throughput. 50 measurements were made for eachchoice of input rate. The measured result is shown in Figure 6.42. We plotted means of thethroughput with their errors (rms) against the input rate. At lower input rate below thetheoretical limit of Gigabit Ethernet, 949.28 Mbps, the output rate is proportional to the
input rate. The theoretical limit of Gigabit Ethernet is calculated by:

176

CHAPTER 6. SENSOR

CMC Test Board

DAQ PCGESiCA
Test Start Command

FPGA
Readout
Controller
Frame
MemoryController

16-bit Test Data +2-bit Control Signals

Checker

Gigabit
Ethernet
PHY
DeviceTCP Client
1000BASE-T

CMC
Connector

VME
Crate

SiTCP

DDR2 SDRAM
SO-DIMM

5V
TCP Server
12V

Regulator

3.3V

Test Data Generator (GESiCA)
FPGA
DIN96
Connector

Test Data
Generation Circuit

CMC
Connector

Figure 6.41: Schematic of the throughput measurement

MSS
IPG + HMAC + HIP + HTCP + MSS= 949.28Mbps

Tmax = 1Gbps ×

(6.2)
(6.3)

Here, M SS is the maximum segment size of TCP, 1460 Bytes; IP G is the inter-packet gap
deﬁned by the speciﬁcation IEEE802.3, 12 Bytes; HM AC is the overhead length of Ethernet,
26 Bytes; HIP is the length of the IP header, 20 Bytes; and HT CP is the length of the TCP
header, 20 Bytes. The errors are so small that they cannot be seen on the graph, whichindicates that the throughput of GESiCA is stable. Throughout this test, we observed noerror in the checker and we conﬁrmed GESiCA has suﬃcient data transfer accuracy.
6.2.3.4

Noise estimation of GESiCA

Jitter measurement of GESiCA
Setup:
To estimate noise from jitter, we directly measured jitter of the readout clock from GESiCA.
Figure 6.43 shows a setup of our jitter measurement. A digital sampling oscilloscope, Tektronics DPO7254, with 40 GHz sampling speed and 250 fs timing accuracy and the software
TDSJIT3 installed in the oscilloscope are used in order to analyze the jitter. We sampled thereadout clock with an active probe on the CMC connector. The jitter for various integration
periods was measured. The integration period is varied by conﬁguring various readout clock
patterns by the DAQ PC.

6.2. FRONT-END ELECTRONICS (FEE) AND BACK-END ELECTRONICS (BEE) 177

Figure 6.42: Meausrement of throughout

Test Readout ClockDAQ PCGESiCA
FPGA
Oscilloscope

Active Probe

Readout
Controller
Frame
MemoryController(Tektronics DPO 7254)

CMC
Connector

SiTCP

Gigabit
Ethernet
PHY
Device1000BASE-T
DDR2 SDRAM
SO-DIMM

3.3V
5V

Power
SupplyFigure 6.43: Schematic of the jitter measurement

178

CHAPTER 6. SENSOR

Figure 6.44: Measurementof jitter

ResultFigure 6.44 shows the measured result. The Y-axis is the ratio of the rms jitter to theintegration time, in parts per million (ppm). The X-axis is the integration time. Errors inthese plots come from the measurement accuracy of TDSJIT3 shown in its data sheet. The
typical integration time is expected to be 2 µs which corresponds to the gain of 1 in the
HSC system. In this case, the jitter of the signal from GESiCA is 10 ppm. To neglect theinﬂuence of readout noise from BEE, equivalent noise charge from the jitter should be lower
than that from a CCD (4 e) around a bias voltage. A standard value for the bias voltage
is equivalent to 1000 ADU. This means that voltages of bias and signal are 1000 ADU. The
noise voltage of readout system, N , is described by
N=

q
2
2 δT
Vsignal+ Vbias.
T(6.4)

Here, Vsignal and Vbias are average voltages of charge signal of the CCD and the bias, respectively. T and δT are an average of the integration time and ﬂuctuation of the integration
time, respectively. The ﬂuctuation is called jitter and is the cause of readout noise. We canestimate readout noise from GESiCA using above equation and measured jitter, 10 ppm, as
follows.
N∼p10002 [ADU] + 10002 [ADU] × (10 × 10−6 ) = 0.014[ADU].

(6.5)

Typical conversion factor is Fc , 3 e/ADU, ∼200,000/216 because the full well capacity of
a CCD is 200,000 e and the ADC is 16bit resolution. Hence the equivalent noise charge
around the bias level is6.2. FRONT-END ELECTRONICS (FEE) AND BACK-END ELECTRONICS (BEE) 179

Nequiv = N × Fc = 0.014 × 3 ∼ 0.04[e].

(6.6)

Thus we concluded that the noise caused by GESiCA is negligible because the noise is much
lower than the noise of the CCD output. Note that the relative accuracy of integration time
1 µs is also about 10 ppm. This means that it is possible to read data from ADC at higher
speed in future updates.
Analog noise measurement of the prototype system
Setup:
Figure 6.45 is the photograph of a prototype readout system which consists of a GESiCA
and front-end electronics for Suprime-Cam. This system consists of three 3U Euro card
size boards (MF2 PWR, MF2 SIG, I/F board), GESiCA, and a backplane. MF2 PWR,
MF2 SIG, and the backplane are designed for FEE of the current Suprime-Cam called
MFront2. A user can control MFront2 by sending commands from a DAQ PC over RS485.
Figure 6.46 shows a schematic of this system. MF2 PWR is a power management board.
MF2 SIG is a signal processing card and processes CCD signals. Two channels are used inthis test. Each channel has one CDS, one 16bit DAC, and one 16 bit ADC. MF2 SIG controls a CDS and an ADC by using readout clocks from GESiCA. Several sets of the digitized
16-bit serial data from the output of each channel are arranged by MF2 SIG into the single
set of 16bit parallel data and sent to GESiCA with a strobe signal and a valid signal. The
range of DAC output (the bias voltage of CCD signal) is from -3 V to 3 V and it can becontrolled by a DAQ PC through MF2 PWR. The input range of the ADC is from 0V to
3V. The CCD signal is integrated by CDS and digitized by the ADC in the MF2 SIG. The
digitized data is transferred to GESiCA, stored in the frame memory temporarily, and thensent to the DAQ PC via Ethernet.
Result:
To obtain characteristics of readout noise of the prototype system, we measured ﬂuctuationof readout values from the MF2 SIG by changing the output voltage of the DAC from 32,768
ADU to 0 ADU with 512 ADU interval (corresponding to the input to the CDS from 0 V to
-3 V in ∼0.047 V decrements). For each voltage, the data processing sequences are repeatedfor 16384 times.
Figure 6.47 shows the measured values. The left and the right Y-axis are the readout noisein ADU and equivalent noise charge respectively. The X-axis is input voltage of the CDS output voltage of the DAC. The solid curves show the measured result. The solid lines show
expected noise calculated from the equation (6.4). .
The result shows that the measured jitter of the readout clock is 14 ppm. The noise levelat the bias level (1000 ADU) is about ∼1.8 e, and the noise come from the jitter is about∼0.06 e. The noise come from the jitter is much lower than the noise of the bias level, (∼1.8e). Hence the noise caused by GESiCA is negligible.
6.2.3.5

Conclusion

We have developed the prototype of GESiCA for the HSC. GESiCA is compact (149 mm
× 79 mm × 12 mm), light weight (∼100 g) and has low power consumption (∼15 W).

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CHAPTER 6. SENSOR

GESiCA + I/F Board

MF2_SIG

MF2_PWR

RS485

Ethernet Connection to the DAQ System
Power for MF2_PWR and MF2_SIG

Figure 6.45: Setup of the analog noise measurement

Command for MFront2 (via RS485)

MF2_PWR

DAQ PCI/F Board

Readout Clock

GESiCA

MF2_SIG
Ch2
DAC

Bias Voltage

CDS

CCD Analog Data (not connected)

Ch1

Ethernet

Command

Readout
Controller
TCP/IP
Processor

ADC

Frame
MemoryDigital Data

Figure 6.46: Schematic of the analog noise measurement6.2. FRONT-END ELECTRONICS (FEE) AND BACK-END ELECTRONICS (BEE) 181

Figure 6.47: Analog Noise Measurement using GESiCA

The important performance factors of GESiCA are throughput and readout noise. As aresult of our performance tests, we conﬁrmed that GESiCA can transfer data through a
Gigabit Ethernet link not only at the speed requirement (0.9 Gbps) but at close to thetheoretical upper limit of 0.949 Gbps. We estimated the noise through the jitter of thereadout clock from GESiCA. The measured jitter is 14 ppm which corresponds to ∼0.6 e,satisfying the requirement for noise. We developed the prototype readout system which
consists of GESiCA and the analog front-end of the Suprime-Cam in order to estimate noiseof the readout system of the HSC. The noise level caused by GESiCA is extremely small.
We concluded our prototype satisﬁed the required speciﬁcation.

6.2.4

Schedule

We plan to make a prototype of BEE and FEE in 2009, and build test equipments for themby the middle of 2009. We would like to establish the procedure of their mass production
and screening tests, and complete the ﬁrst production by early in 2010.
A vacuum chamber is needed to evaluate the heat exhausting eﬃciency of the aluminum
core PCB. We have to verify that the temperature rising could be kept within an acceptable
range. The chamber is also valuable for screening tests such as a burn-in test of the FEE
board after the beginning of the mass production. It is important to reduce initial failures
to avoid a repetition of the replacement work after the installation into the dewar as small
as possible.
We also have to make a ’mosaic CCD simulator’ to check FEE working properly, whichtakes care of 4 CCDs. It has 16 signal outputs, and several 10 channels of voltage monitors

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CHAPTER 6. SENSOR

to check the clock and bias voltages come from an FEE board.
Furthermore, we need to develop some software to execute the screening test automaticallyto check several 10 FEE boards, and collaborate with software groups to integrate dataacquisition system with our hardware for both of science CCDs and guide CCDs.

Bibliography[1] Miyazaki, S. et al. 1998, SPIE, 3355, 363[2] Kamata, Y. et al. 2004, SPIE, 5499, 210[3] Kamata, Y. et al. 2006, SPIE, 6276, 52[4] Kamata, Y. et al. 2008, SPIE, 7021, 52[5] Miyazaki, S. et al. 2002, PASJ, 54, 833[6] Kawanomoto, S. 2009, HSC Report, HSC20090217 KawanomotoS CCDdepth[7] Janesick, J. 2001 “Scientiﬁc Charge-Coupled Devices” p.255[8] Gunn, J. E. et al., 1998, AJ, 116, 3040
[9] Miyatake, H. et al., 2008, IEEE Nuclear Science Symposium, N02-25, 737[10] Miyatake, H., 2009, University of Tokyo, Master thesis[11] Nakaya, H. et al., 2004, SPIE, 5492, 852[12] Nakaya, H. et al., 2006, PASP, 118, 478[13] Nakaya, H. et al., 2006, SPIE, 6269, 127[14] Nakaya, H. et al., 2008, SPIE, 7014, 162[15] Uchida, T., 2008, IEEE Transactions on Nuclear Science, 55, 1631
[16] Uchida, T. et al., 2008, SPIE, 7014, 238

183

Chapter 7
FilterIn this chapter, we present speciﬁcations and implementations of broad band ﬁlters for HSC.
Furthermore, we introduce current status of prototyping of the ﬁlters.

7.1

Broad band ﬁlters for HSC

We are planning to provide the HSC with broad band ﬁlter set equivalent to that for theSuprime-Cam. The Suprime-Cam has two photometric band systems for broad-band ﬁlters,
the Johnson-Morgan-Cousins system and the SDSS system[1].We selected the SDSS system
for the photometric system for the HSC from these two systems. In addition to this, we areplanning to produce the Y band ﬁlter[2].

7.1.1

Requirements

The broad band ﬁlters consist of multi-layer interference coatings deposited on plane-parallel
synthetic silica substrates. We present detailed speciﬁcations for the broad band ﬁlters inthis section.

7.1.1.1

Physical properties

Since the HSC has three times larger focal plane diameter than the Suprime-Cam, it is inevitable that the requirement for physical size of ﬁlters becomes large in the same proportion.
So a ﬁlter should cover a circle 500 mm in diameter at the focal plane (see Camera chapter).
Since the ﬁlter will be placed at about 70 mm in front of the CCDs and the F ratio of WFC
optics is about 2, the clear aperture of the ﬁlters should be larger than about 570 mm. We
assign 11mm margin at the both side of clear aperture and adopt 582 mm for the diameterof clear aperture. In addition to this, there are 9 mm blank space to hold ﬁlter substrates
at the edge of clear aperture. So we adopt 600 mm as the diameter of substrates. And, theback-focus of the WFC is only about 180 mm, the thickness of ﬁlter frame is very limited.
Because a ﬁlter by handled by an automatic ﬁlter exchange system, there is a tight limit
in weight of ﬁlter and ﬁlter frame to make the system as reliable as possible. We tabulaterequirements for these physical properties in Table 7.1.
185

186

CHAPTER 7. FILTER

Item
Diameter
Clear aperture
Thickness
WeightRequirements
600 mm
582 mm< 42 mm (including ﬁlter frame)< 25 kg (including ﬁlter frame)

Table 7.1: Requirements for physical properties for a ﬁlter.

7.1.1.2

Optical properties

We deﬁne a sub-aperture (SA) as a cross section of converging beam bisected by the ﬁlter
surface, and the diameter of the SA is about 35mm. Optical properties of a ﬁlter should bedetermined not to aﬀect image quality in any SA in a ﬁlter surface. And, since a ﬁlter is inconverging beam, error of thickness brings spherical aberration to the image. In addition,
because a bat parallelism of both surfaces of a ﬁlter brings a tilt of the focal plane, theparallelism between surfaces of substrate should be provided. We tabulate requirements for
optical properties in Table 7.2.
Item
Requirements
Wavefront error 100 nm RMS over SA
Parallelism
40 00
Thickness error < 0.5 mm
Table 7.2: Requirements for optical properties for a ﬁlter.

7.1.1.3

Spectral properties

Most of Suprime-Cam ﬁlters adopt color glass for the one end of the band pass but there isno such large color glass that can be used for HSC. We however try to mimic the SuprimeCam’s transmission characteristic as much as we can using interference coating only. A
typical values of characteristics for ﬁlters of the Suprime-Cam are shown in Table 7.3. We
Item
Cut-oﬀ wavelength error
Variance across clear aperture
Peak transmission
Out of band leak
Ripple of transmission

Actual< 7 nm∼ 2-4 nm
> 95 %< 0.1 %
∼5%

Table 7.3: Typical spectral properties for the Suprime-Cam ﬁlters.

have to relax the speciﬁcation for “out of band leak” compared with the Suprime-Cam

7.1. BROAD BAND FILTERS FOR HSC

187

value because of the lack of color glass ﬁlters. Since the WFC optics has no transmission
at the wavelength below 350 nm (see WFC chapter) and the detectors have no sensitivity
at the wavelength above 1050 nm, we chose the wavelength range of 350-1050 nm as therequirements for the blocking of ﬁlters. We tabulate requirements for optical properties inTable 7.4.
Item
Cut-on/-oﬀ wavelength error
Variance across clear aperture
Peak transmittanceOut of band leak
Ripple of transmission
Blocking wavelength range

Requirements< 10 nm< 2 % of central wavelength> 90 %< 0.5 %
<5%350 − 1050 nm

Table 7.4: Requirements for spectral properties for ﬁlter set.

7.1.1.4

Summary of Speciﬁcations

We summarize the speciﬁcations for a ﬁlter of HSC in Table 7.5. The measured values of
the Suprime-Cam ﬁlters are also shown for reference.
Item
Substrate
Diameter
Clear aperture (CA)Thickness
Thickness error
Parallelism
Diameter of SA
Wavefront error in SA
Peak transmittanceOut of band leakCut-on/-oﬀ wavelength error
Uniformity(wavelength)
Uniformity(transmission)
Ripple of transmission
Blocking wavelength range

Speciﬁcations
Synthetic silica
600 mm
582 mm< 42 mm< 0.5 mm
< 40 00
35 mm
100 nmRMS> 90 %< 0.5 %
<2%
<2%
<2%
<5%
350-1050 nm

Goal

30 nmRMS> 95 %< 0.1 %
< 0.5 %
< 0.5 %
< 0.5 %
<2%
300-1200 nm

Suprime-Cam
BK7 & color glass
205 × 170 mm
192 × 158 mm
20 mm (w/ holder)< 0.2 mm
<1030 mm
> 95 %< 0.1 %
∼ 0.6-1.0 %∼ 0.4 %
∼1%
∼5%

Table 7.5: Requirements for spectral properties for ﬁlter set.

7.1.1.5

Expected transmission curves

We present wavelength coverage of the HSC broad band ﬁlter set in Table 7.6. For comprehension, we present expected transmission curves of the broad band ﬁlter set in Figure 7.1.

188

CHAPTER 7. FILTER

Band
g0
r0
i0
z0
z0 (optional)
Y

Wavelength [nm]400-550550-695695-845845-960845-930960-1070

Table 7.6: Wavelength coverage of the HSC broad band ﬁlter set.

The “optional” z’ ﬁlter is also considered. This optional band may be arranged to avoid thewater vapor absorption at 950 nm which is variable and annoying.

7.1.2

Implementation

Interference ﬁlters used to have somewhat poorer performance than color glass ﬁlters for outof band blocking, however, some promising prototyping is now carrying out. Since there areseveral types of deposition method, we try to test normal vacuum deposition and sputter
deposition in our prototyping.
Normal vacuum deposition method can produce uniform coating on large substrate , however, there are some problems to be discussed. First, owing to the limitation of the thicknessof deposition layer, at least ﬁve coating surfaces are required to block wide wavelength rangeof out-of-band and thus three thin substrates are essential. Multi-substrate ﬁlter will cause
a ghost image around bright stars and thin substrate must be treated with great care inassembling and handling. Request from the WFC optics and the ﬁlter exchanger, thicknessof a substrate is limited to only 7 mm. And secondly, the normal vacuum deposition coating layer soaks up water vapor molecules and changes its refraction indexes, therefore theproperties of ﬁlter vary with environmental humidity variation.
Since the sputtering process can build dense coating layer on a substrate, resistance tothe environment ought to be better than those of normal vacuum deposition coatings. If
sputtering process meets our speciﬁcations, we can reduce the number of substrates to one
because thicker deposition layer is feasible. This will reduce the strength of ghost images.

7.1.3

Prototyping at OCJ

We placed an order for a prototype broad band ﬁlter to OCJ (Optical Coatings Japan) inorder to verify that a band pass ﬁlter consists of interference multi-layer coating is acceptableto our purpose. As OCJ use normal vacuum deposition method to make interference ﬁlters,
we tried to test whether normal vacuum deposition coating meets our requirements or not.
To save money and time for the prototype, we did not use synthetic silica substrate in theﬁrst prototyping and the blocking layers were omitted from the ﬁrst prototyping of broadband ﬁlter. We present the speciﬁcations and the results for the ﬁrst prototyping in Table
7.7. The transmission curves and a photograph of the prototype ﬁlter are also presented in
Figure 7.2 and 7.3.

7.1. BROAD BAND FILTERS FOR HSC

189

100

Transmission [%]

80

g

r

i

z

Y

900

1000

60

40

20

0
300

400

500

600

700

800

1100

1200

Wavelength [nm]Figure 7.1: Expected transmission curves for HSC ﬁlter set.

The result shows the prototype has better transmittance and better uniformity than expected.

7.1.4

Prototyping at BARR

In a prototyping for narrow band ﬁlter promoted by T. Hayashino(Tohoku Univ.), it turned
out that the sputter deposition coat is necessary because it requires more layers than broad
band’s. We ordered a prototype narrow band ﬁlter to BARR (BARR associates, inc.)through the representation in Japan, FUJITOK corporation. BARR has a large chamber for sputtering. We present the speciﬁcations and the results for the ﬁrst prototyping of
narrow band ﬁlter in Table 7.8. The transmission curves of the prototype narrow band ﬁlter
are also presented in Figure 7.4. We are satisﬁed with the result of this prototype because
it fulﬁlls the speciﬁcation for uniformity and has suﬃciently sharp cut-oﬀ edge and has very
good transmittance.
The result of the ﬁrst prototyping of narrow band ﬁlter shows us a great possibility ofthe sputter deposition in making broad band ﬁlters. However, BARR does not guarantee
speciﬁcations in the outer edge of the clear aperture of r < 250mm which we do not fully
understand why. In fact, the result of prototyping of the narrow band ﬁlter satisﬁed thespeciﬁcation even at the edge of ﬁlter (see Figure 7.5). So we are now contracting BARR
about the speciﬁcation for the diameter in more detail.
Here we introduce the ﬁlter measuring system. To measure the ﬁlter in F/2 converging

190

CHAPTER 7. FILTER

Item
Substrate
Diameter
Clear aperture
Thickness
Thickness error
Parallelism
Wavefront error
Cut-on wavelength
Cut-oﬀ wavelength
Peak transmittanceOut of band leakSpeciﬁcationsR
TEMPAX Float°
600 mm
580 mm
10 mm
N/A
N/A
N/A
550 nm695 nm> 90 %<1%

Cut-on/-oﬀ wavelength error< 1.5 %

Uniformity(wavelength)
Uniformity(transmission)
Ripple of transmission
Blocking wavelength range

<2%
<2%
<5%
465-810 nm

Results
600 mm
580 mm
10 mm

545 nm692 nm
94 % (avg in 565-680 nm)
0.4 % (avg in 465-525 nm)
0.2 % (avg in 730-810 nm)
0.9 % @ 550 nm
0.4 % @ 695 nm
0.5 %
1%
2%

Table 7.7: Speciﬁcations and results for the ﬁrst prototyping of broad band ﬁlter at OCJ.

light, we made ﬁlter measuring system in NAOJ. The system has two degree of freedom (r, θ)and is attached to a spectrophotometer (Shimadzu UV3100PC) in Opto-shop of ATC/NAOJ.
In Figure 7.6, we present a photograph of the measuring system. In this ﬁgure, prototype
narrow band ﬁlter is shown too.

7.1.5

Filter handling

Since a substrate has only 7mm thickness is the case of normal vacuum deposition, special
care will be required to handle it. Once it assembled in ﬁlter frame, there is not much riskof destruction. Proper assembling sequence should be adopted.

7.1.6

Durability

Normal deposition coating easily absorb moisture in through the surface because the kinematic energy of deposition particle is small and therefore the deposition layer tends to beporous. Because of changes of refraction index, cut-oﬀ wavelength is sensitive to environmental humidity. Study of this eﬀect on the test piece of prototype ﬁlter has been done.
We present the result here. We set the environment temperature at 293 K, and set relative
humidity at 30 % and 85 %. Between these situations, the prototype ﬁlter shows the wavelength shift of cut-oﬀ edge of the band-pass and the amount of the shift, ∆λ, is 1.5 nm and3.5 nm for leading and trailing edge, respectively. Since the saturation speciﬁc humidity is
23.4 hPa at 293 K and 6.11 hPa at 273 K, the wavelength shift of the ﬁlters at the summit isexpected about a thirds of this study, and furthermore, as the band-pass interference surface

7.1. BROAD BAND FILTERS FOR HSC

191

100

Center
5cm
10cm
15cm
20cm
25cm

90
80

Transmittance [%]

70
60
50
40
30
20
10
0
400

450

500

550

600

650

700

750

800

Wavelength [nm]Figure 7.2: Transmittance curves for prototype broad band ﬁlter.

will be mounted inside the ﬁlter frame (see next section), the wavelength shift will be muchsmaller than the calculated value. For detail, see the corresponding report [4].
Here we present expected maximum error on photometry in the wavelength shift of prototype broad band ﬁlter. In the case of 10000 K black-body radiation and the relative humidity
changes from 0 % into 100 % at the summit, there is a diﬀerence about +0.005 mag in photometry. For 3000 K black-body radiation, the photometric change turns to −0.006 magin the same situation of humidity. This photometric error can be suppressed by appropriate assembly of ﬁlter substrates, however, normal vacuum deposition type ﬁlter should bemeasured its spectral properties periodically.

7.1.7

Assembly of three substrates

To isolate the band-pass surface from environment, we stack three substrates in proper order.
The schematic view of ﬁlter stack is shown in Figure 7.7.

7.1.8

Filter glass holder assembly

The ﬁgure of planned ﬁlter glass holder assembly is shown in Figure 7.8. In order to manage
atmospheric pressure change with transportation from the sea level to the summit, air relief
screws are located at support ring. These screws will be shut after arriving the summit.

192

CHAPTER 7. FILTER

Figure 7.3: Photograph of the prototype broad band ﬁlter.

Concept of ﬁlter frame and ﬁlter glass holder is also shown in Figure 7.9. And, we are
discussing assembling sequence of ﬁlter glass holder. Concept of substrate stacking is shownin Figure 7.10. To tighten screws from bottom side of substrate stack, we plan to prepare
a jig. We are now discussing how to compensate the manufacturing tolerance of thickness
and diameter of ﬁlter substrates. Thin spacer and small wedge can be acceptable to holdﬁlter substrates.

7.1.9

Ghost images

In the three substrate (and not glued) ﬁlter, ghost images around bright stars may be concern
because there are a lot of plane parallel surfaces. We investigated this issue and obtained
results. A schematic view of this situation can be seen in Figure 7.11. We estimated the
proportion of the aﬀected area to the FOV of HSC with the same criteria in the reduction
procedure of Suprime-Cam images. We assume that the lens ﬂare (see Figure 7.11) has thesame strength ratio to the central star and radius of it on HSC. Under this condition, we
found that the percentage of aﬀected area does not change much between the Suprime-Cam
case and the HSC case. To be precise, we found that a percentage of masked area is 2.11% in
the case of HSC and 2.26% in Suprime-Cam. We interpret this result as the following. The7.1. BROAD BAND FILTERS FOR HSC

193

Item
Substrate
Diameter
Clear aperture
Thickness
Thickness error
Parallelism
Wavefront error
Central wavelength (CW)
Band-pass width (BW)CW error

Speciﬁcations
Results
Schott B270
600 mm
600 mm
580 mm
580 mm
15 mm
15 mm
N/A
N/A
N/A
600 nm
600 nm8 nm
8 nm< 1.7 nm(r < 200 mm) 1.7 nm< 3 nm(r < 250 mm)
3 nm
BW error< 1.5 nm(r < 200 mm) 0.6 nm< 2 nm(r < 250 mm)
0.8 nmBlocking wavelength range
N/A
Table 7.8: Speciﬁcations and results for the narrow band ﬁlter prototyping at BARR.

narrow air space and thin substrate make rather small ghost image. So the most prominent
ghost in HSC image can be hidden behind the mask already placed to hide the lens ﬂare.
Therefore we conclude that the eﬀect of ghost image is not severe in HSC. For detail, see
the corresponding report [3].

7.1.10

Prototyping and development status / schedule

We summarize the schedule of prototyping and development of ﬁlters in Table 7.9 and 7.10
for the cases of OCJ and BARR, respectively. We have a choice of two method of deposition
but this schedule does not change drastically in either case because they have same delivery
date of six month for a ﬁlter.
At the date of starting system assembly (the third quarter of 2010), we will obtain at least
three ﬁlters, and the complete ﬁlter set that consists of ﬁve ﬁlters will be completed by thethird quarter of 2011.

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CHAPTER 7. FILTER

Date-2008/12-2009/01
2009/01-

Eventsprototype #1 completed
environmental test
prototype #2
Study for product version (three substrates)
If prototype #2 has good properties, we willuse it as a component of the product version-2009/03 ﬁrst substrate of prototype #2 will complete-2009/11? third substrate of prototype #2 will complete
2009/12- test prototype #2 & start next one
If the result of this test is acceptable for product use,the ﬁrst product version of ﬁlters will completion here (r’ band).
-2010/3Q completion of product version #2 and #3-2011/3Q completion of product version #4 and #5(completion of full ﬁlter set)
Table 7.9: Development schedule (OCJ case).

Date
2008/-10
2009/04-

Eventsprototype #1 completed
start the ﬁrst sputter deposition broad band ﬁlter(product version)
2009/11? the ﬁrst sputter deposition broad band ﬁlter will complete-2010/3Q completion of product version #2 and #3-2011/3Q completion of product version #4 and #5(completion of full ﬁlter set)
Table 7.10: Development schedule (BARR case).

7.1. BROAD BAND FILTERS FOR HSC

195

100

∆CW~3nm8060BW~8nm
40

20

0
590

595

600

605

Figure 7.4: Transmittance curves for prototype narrow band ﬁlter.

610

196

CHAPTER 7. FILTER

Figure 7.5: CW variation with radial distance for prototype narrow band ﬁlter.

7.1. BROAD BAND FILTERS FOR HSC

Figure 7.6: Figure of ﬁlter measuring system.

197

198

CHAPTER 7. FILTER

Cut filter 3
Cut filter 1
Substrates
Air gaps

Short pass filter
Long pass filter
Cut filter 2
Cut filter 4

Figure 7.7: Schematic view of ﬁlter substrates in ﬁlter frame.

PTFE Pads

Glass holder

Air relief
screwsFilter glasses

Support ring

Figure 7.8: Figure of glass holder.

7.1. BROAD BAND FILTERS FOR HSC

Support ring
- 580mm aperture- Made of Aluminum
Air relief screw- Purge enclosed air
when pressure changes

Glass holder- 580mm aperture- Compatible with conic coupling
which is proposed by ASIAA/ASRD- Made of Aluminum- Weight 2.2kg
Filter frame- Provided by ASIAA/ASRD

Gap between filter frame and glass holder- To avoid stress caused bydeformation of the filter frame
PTFE Pads- 4 ring pads and 1 lateral pad- Groove for air relief

Figure 7.9: Figure of glass holder.

199

200

CHAPTER 7. FILTER

Thickness gauge
Jig

Figure 7.10: Figure of glass holder.

7.1. BROAD BAND FILTERS FOR HSC

Figure 7.11: Schematic view of ﬁlter ghost and ﬂare around bright star.

201

Bibliography[1] Fukugita, M., Ichikawa, T., Gunn, J. E., Doi, M., Shimasaku, K., Schneider, D. P. 1996,
AJ., 111, 1748
[2] Hillenbrand, L. A., Foster, J. B., Persson, S. E., Matthews, K. 2002, PASP., 114, 708[3] Kawanomoto, S. 2009, HSC Report, HSC20090126 KawanomotoS FilterHumidity[4] Kawanomoto, S. 2009, HSC Report, HSC20090209 KawanomotoS FilterGhostNo4

203

Chapter 8
Telescope Interface
For mechanical and operational interface between the Subaru telescope and HSC, we needsome modiﬁcations of the current systems of the telescope to realize HSC. We also have todevelop the HSC systems meeting speciﬁcations of the telescope systems. In this section,
we describe necessary modiﬁcations of mechanical and operational interface between thetelescope and HSC. We also discuss possible collisions when transferring HSC to the top
unit and ideas for avoiding these accidents.

8.1

Telescope-HSC System Conﬁguration

The telescope-HSC system conﬁguration diagram is shown in Figure 8.1. The rightmost box
is HSC and the left side is operation systems which telscope operators and observes can see
and operate. Colored boxes in the ﬁgure are what we need to bulid newly (green; NAOJ,
pink; MELCO, and light blue; CANON) and modify (orange; MELCO). Green (NAOJ) andred (MELCO) lines are cables newly added.
We ﬁrst describe major changes related to the operation of other instruments and foci. For
HSC, interface of FRCU will be a newly-built control system (VME/PCI), which is diﬀerent
from that of the existing control system (VME/VME), and we need to reconnect instrument
rotator encoder cables everytime we attach HSC or FMOS to the top ring. For secondary
mirror units, auto-connectors will be removed because the inside of the inner hub is trimmed(section 8.3.1).
Next, we describe relatively minor changes which do not aﬀect other instruments. TWS1
and TSC are interfaces which telescope operators and observers can see and operate directly
and interface to the HSC operations should be added for them. Modiﬁcations of MLP1 and
MLP2, which are connected to four HSC control systems (SCAGCU, SCPACU, SHAGCU,
and SHCU), are required because we made new systems for AG and SH as described in
section 8.2.1 and section 8.2.2, respectively. Three pink boxes at the center of the ﬁgure
are control systems for the HSC instrument rotator, hexapod, and cable switching system(section 8.3.3).
In summary, there are no special or serious impacts on other instruments, exchange of
instruments and secondary mirrors. More details for the telescope-HSC system conﬁguration
are shown in the document [1].

205Figure 8.1: Telescope-HSC system conﬁguration diagram. See the text for details.

206
CHAPTER 8. TELESCOPE INTERFACE

8.2. DESIGNS FOR AUTOGUIDER/SHACK-HARTMANN/AUTOFOCUS

8.2

207

Designs for Autoguider/Shack-Hartmann/Autofocus

Autoguider (AG), Shack-Hartmann (SH), and autofocus (AF) systems are necessary to maintain good image quality.
AG and SH systems of the current prime focus are installed on the probe between SuprimeCam and the prime focus corrector. All the optics and mechanics are involved in the probe.
However, the tight space constraint does not allow us to adopt the same systems for HSC.
We need to introduce modiﬁed AG and SH systems as below. Conceptual designs for the
AG and SH systems are shown in our SPIE paper [2].
AF system will be newly introduced. The AF system monitors shapes of objects on CCDs
set with vertical oﬀsets at edges of the ﬁeld-of-view to adjust the position of the HSC unit.

8.2.1

Autoguider

Open tracking accuracy of the Subaru telescope is 0.2 arcsec rms per 10 minutes [3]. Typicalseeing size in optical wavelengths at Mauna Kea is subarcsecond and the open tracking candegrade image quality for observations in typical exposure times. More accurate tracking
can be achieved by an AG system which measures positions of a bright AG star continuously
and sends tracking errors to the telescope. The AG system is now installed on the probe
for the current prime focus which is used for Suprime-Cam observations1 . This AG system
is used for science observations using the Suprime-Cam CCDs as well as pointing analysis(PA) and mirror analysis (MA) using the AG CCD on the probe. However, we can not
introduce exactly the same system because we have little space between the HSC focal plane
and wide-ﬁeld corrector.
Conceptual Design
We have designed to use some of the science CCDs at the edges of the ﬁeld-of-view, wherethe vignetting is ∼ 25%, for the HSC AG system. Photons of AG stars must pass throughbandpass ﬁlters used then. The current AG systems have a ﬁxed ﬁlter which photons of∼ 650−810 nm can pass through. This wavelength is selected in terms of the high throughput
and number of bright stars. In order to get similar signal-to-noise ratios for the HSC AG,
it is expected that we need longer exposure in short or long wavelengths where the HSC
throughput is not so high. We have to seriously consider how long exposure times we needand how many bright stars there are.
We have calculated limiting magnituides for all the broad and narrow bands consideredfor HSC survey under any bad weather conditions (low atmospheric transparency and highsky background, see Table 8.1). From the error budget of image quality for AG (section 3),
autoguiding accuracy required is 0.14 pixel rms and this requires signal-to-noise (S/N ) ratioof 50 based on our simulations which depends on seeing size. We here adopt the rms value
when seeing is a typical value, 0.8 arcsec. We found that 3-second exposure will give us
suﬃcient (∼ 50) S/N ratios for 17th magnitude in grizY bands and 15th magnituide in
narrow bands. In typical observing conditions, we take 1-second exposure. We have estimated the cumulative number density of such bright stars > 100 deg−2 based on the Tycho-2
catalog and Sloan Digital Sky Survey (SDSS) database (Data Release 7, [4]) and an areaof about two 2k×4k CCDs is necessary if we guarantee to ﬁnd at least 1 autoguide star inany direction (even for the north Galactic pole). Therefore, we plan to assign four CCDs at
1

Another prime focus instrument, FMOS, has its own AG systems.

208

CHAPTER 8. TELESCOPE INTERFACE

the edges of FoV for AG as shown in Figure 8.2. We have also conﬁrmed that we can reallyﬁnd AG stars of suﬃcient brightness when we point the telescope to some of directions. We
note here that FDCCDs must be adopted for the AG CCDs because high sensitivity in longwavelength is also required for the autoguide system.

Table 8.1: Limiting magnitudes (S/N = 50) for the HSC AG system. texp is an exposure
time in second. The worst condition considered here is 1.0 magnitude extinction and 10
times brighter sky background.
texp
sec
1
3
1
3

u
18.4
19.6
16.9
17.8

g
20.6
21.6
18.6
19.3

broad band
r
i
20.7
20.5
21.6
21.4
18.4
18.2
19.1
18.9

z
19.6
20.4
17.2
17.9

Y
19.0
19.7
16.5
17.1

387
16.3
17.5
14.8
16.0

526
17.5
18.7
16.0
17.1

narrow band (NB)
717
816
921
17.6 17.6 17.5
18.8 18.7 18.5
16.1 16.0 15.6
17.1 16.9 16.3

973
17.1
18.2
15.2
15.9

101
16.5
17.6
14.7
15.4

condition
typicaltypical
worst
worst

Figure 8.2: Locations of the HSC AG CCDs. Dashed lines indicate ﬁeld-of-views of 1.5
degree diameter and 1.6 degree diameter, respectively.

Parameters necessary to enable autoguiding such as telescope pointing error (∆AZ, ∆EL)
based on measurements of diﬀerence of AG star positions (∆x, ∆y) transferred from SCAGCU
to MLP1 through the HSC SCAGCU, which is the same sequence as the current AG system.
One cycle is roughly 6-second loop including the CCD partial readout. We are now assuming
the readout region is ∼ 30 × 30 [arcsec2 ]. AG procedure in real observations is shown in
Figure 8.3. We can not close the HSC shutter during a science exposure and then photons
from sky come into the AG CCDs during the AG CCD readout. We have investigated the8.2. DESIGNS FOR AUTOGUIDER/SHACK-HARTMANN/AUTOFOCUS

209

eﬀect of electron ﬂows caused by such photons and found that they are not a problem fordeterminations of AG star positions.
More details can be found in the HSC Report documents [5].

Figure 8.3: Observation sequence with the HSC AG system.

8.2.2

Shack-Hartmann

The mirror analysis (MA) procedure determines powers of actuators for the primary mirror
as a function of telescope elevation. Shack-Hartmann (SH) system on the prime focus probe
is now used for the MA procedure. SH system for HSC should have the same ability as thecurrent system in terms of a dynamic range and resolution of measurements. However, theexactly same SH system as the current prime focus can not be realized because of tight spaceconstraint as well as the AG system.
Conceptual Design
Required measurement accuracy for wavefront errors for the SH system is 0.02 arcsec, whichis the same for that for the current SH system [6]. Position oﬀsets between reference lightimages and stellar images formed by an microlens array (MLA) are necessary to calculate
Zernike coeﬃcients. Diﬀerence of aberration between reference light and stellar light due todiﬀerent paths should be less than 0.01 arcsec.
We have designed that the SH unit will be installed as a ”SH ﬁlter”. The height of the SH
ﬁlter will be 42 mm and can be installed in the HSC ﬁlter unit. The cartoon picture anddesigned picture are shown in Figure 8.4 and Figure 8.5, respectively. All the optics for theSH system except for the SH CCD are involved in the ﬁlter and stellar images divided by

210

CHAPTER 8. TELESCOPE INTERFACE

the MLA are formed on the science CCD. The SH ﬁlter has its own autoguide system using
the WATEC commercially-produced high-sensitivity camera, WAT-120N2 .
Light from the primary mirror ﬁrst is reﬂected at the pick-up mirror and once focused at
the diaphragm. The collimator lens unit makes it collimated and the collimated beam enters
the half mirror. The collimator lens unit is carefully designed to meet the requirement above
within 3 arcsec diameter ﬁeld-of-view. The reﬂected light passes through a ﬁlter transparent
in 630-690 nm and the next lens makes it a focused image on the AG CCD which is operated
by SHAGCU. The transmitting light passes through a ﬁlter transparent in 500-570 nm andenter the MLA (16 × 16 1.5 × 1.5 mm2 microlens elements) with a focal length of 100 mm.
The light passing through the MLA is reﬂected by the 45-degree mirror and the images areformed on the science CCD. Position oﬀsets between reference lights and stellar lights aresent from the HSC SH control system SHCU to MLP2 to calculate Zernike coeﬃcients.
Based on the experience of the telescope scientist and operators, 30 arcsec diameter ﬁeldof-view is neccesary to get an image of the target star before the PA is done if we conduct
PA using the AG camera in the SH ﬁlter. On the other hand, we need to avoid light fromother stars in the ﬁeld-of-view to come into the SH optics during the MA procedure. In
this case, we need to limit the ﬁeld-of-view ∼ 3 arcsec diameter. Therefore, we introduce amechanical exchange unit by which we can replace diaphragms of 400 µm diameter (3 arcsec
case) and 3000 µm diamter (30 arcsec case). PA can be also done using autoguide CCD
in the SH ﬁlter as well as the science CCD at the center of the ﬁeld-of-view. We are now
examining which CCD we use for the PA procedure.
We have started fabrication of a prototype of the SH ﬁlter. Almost all the components
should be set with an accuracy of ∼ 50 µm except for some parts requiring 10 µm accuracy.
There are also other parts with alignment mechanisms which can adjust their positions with
10 µm accuracy.
More details for determination of focal length of the MLA, for example, can be found in
the HSC Report document [7].

8.2.3Autofocus

Autofocus (AF) system is not currently introduced for Suprime-Cam. However, realtime
focusing would give us stably good image quality and save focusing time which is now
typically done a few times during a night. We put CCDs at diﬀerent height (Figure 8.6)
along optical axis at edges of the ﬁeld-of-view (Figure 8.8). The height diﬀerence from thescience CCDs is ±200 µm. For Suprime-Cam focusing images taken on October 25th, 2008,
we measured full-width-half-maxima (FWHMs) of several stars at diﬀerent focus positions
and found the best focus position automatically (Figure 8.7). We need more examinations
such as similar tests using simulated HSC images for which vignetting and real abberation
are taken into account and best parameterization for accurate determination of the best
focus position.
Images for AF are read out at the same time as science CCDs are read out and we calculate
the best focus position to give a feedback to the HSC unit attitude. In typical exposure time,
the number of objects in each AF CCD is a few tens in typical weather conditions.
2

http://www.watec.co.jp/english/bw/wat 120n+.html

8.2. DESIGNS FOR AUTOGUIDER/SHACK-HARTMANN/AUTOFOCUS

211

Figure 8.4: A cartoon picture of the HSC SH ﬁlter from a edge-on view.

Figure 8.5: Optics layout of the HSC SH ﬁlter based on the mechanical design. We notethat light from the primary mirror comes from the downside in this ﬁgure and inverse of
Figure 8.4.

212

CHAPTER 8. TELESCOPE INTERFACE

Figure 8.6: A cartoon picture of the HSC AF concept. AF CCDs are put with oﬀsets of±200 µm from the science CCDs.

Figure 8.7: FWHMs of stars in the focusing test images of Suprime-Cam and 2ndpolynomials ﬁtted best for each star (colored points and solid lines) and all (black dashedline).

8.3. MECHANICAL MODIFICATION OF TELESCOPE

213

Figure 8.8: Locations of the AF CCDs.

8.3

Mechanical Modiﬁcation of Telescope

In order to realize a wide ﬁeld-of-view of HSC, mechanical interfaces between the telescopeand HSC need to be modiﬁed. Our basic policy is that impacts on observatory operationand performance of other instruments should be minimized.

8.3.1

Inner hubA wider ﬁeld-of-view with small vignetting requires larger corrector lenses. The inner hub
is a component at the top ring, which supports units attached to the prime focus such as
the prime focus unit, secondary mirrors, and HSC. The diameter of the inside of the current
inner hub is 910 mm and we will expand it to 1000 mm as shown in Figure 8.9. However,there is small clearance between the inner hub and HSC when HSC is attached to the prime
focus from the top of the inner hub. The HSC base frame is attached to the inner hub withthe clipping structure (section 8.3.2). We have to take much care of possible accidents which
will be discussed in section 8.5.

8.3.2

Clipping mechanism

At the top of the inner hub, we put the interface of the HSC base frame, and connect these
two structures by the clipping mechanism as shown in Figure 8.10 and Figure 8.11. We needto make M10 screw holes in the inner hub. This clipping structure can be involved inside theinner hub and there are no inﬂuence on other instruments especially for infrared instruments
and adaptive optics.

214

CHAPTER 8. TELESCOPE INTERFACE

Figure 8.9: Modiﬁcation of the inner hub of the top ring. We trim the ﬂange inside of theinner hub to expand the inner diameter from 910 mm to 1000 mm.

Figure 8.10: Clipping mechanism.

8.3. MECHANICAL MODIFICATION OF TELESCOPE

215

Figure 8.11: Clipping procedure.

Details are shown in the document [1].

8.3.3

CablesWe trim the ﬂange inside of the inner hub where auto-connectors for secondary mirrors areattached. There are now two auto-connectors for driving secondary mirrors and tip-tilting
which connect with cables on the front-side spider as shown in Figure 8.12. After the insideof the inner hub is trimmed, these connectors have to be connected manually. The connector
on secondary mirrors are connected with a jumper cable to the prime focus connector on
the rear-side spider (Figure 8.13). The connector for tip-tilting which is used for the IR
secondary mirror is connected with another jumper cable to the same connectors as original
one (Figure 8.14). The laser launch telescope (LLT) is set on the top of the inner hub. Partof the curvic coupling of the LLT which connect to the inner hub has been already detachedand cables can pass through the holes to connect the cables on the spiders (Figure 8.15).
Cables between HSC and the control system on the second ﬂoor of the control building
can be an issue because we may have to use backup wires. Moreover, if we would connect
HSC with the control system with cables exclusively for HSC, extensive constructions forcabling are needed. This is unrealistic. Instead, we put a new cable switching system on thesecond ﬂoor of the control building and share the existing cables between the prime focus
connector box and control building with other prime focus instruments (Figure 8.16). As for
optical ﬁber cables, it is not necessary to modify the AZ/EL cable wrapper because we use
backup cables.
Details are shown in the documents [8, 9].

8.3.4

Telescope Balancer

HSC is heavier than Suprime-Cam and we have to consider whether it is necessary or not
to modify the current telescope balancer (Figure 8.17). Our basic policy is that telescope

216

CHAPTER 8. TELESCOPE INTERFACE

Figure 8.12: Current cable connection when the secondary mirror in use.

Figure 8.13: CS-Opt and NS-Opt secondary mirror in use after the inner hub modiﬁcation.

8.3. MECHANICAL MODIFICATION OF TELESCOPE

217

Figure 8.14: IR secondary mirror without LLT in use after the inner hub modiﬁcation.

Figure 8.15: IR secondary mirror with LLT in use after the inner hub modiﬁcation.

218

CHAPTER 8. TELESCOPE INTERFACE

Figure 8.16: Cable switching system.

modiﬁcation would be as small as possible. Therefore, we set constrains on the HSC weightfrom viewpoints of i) TUE laden weight constraint (TUE can carry less than 3.0 tons), andii) the telescope balancer which must not be modiﬁed signiﬁcantly.

Figure 8.17: Current telescope balancers.

Constraint from the TUE laden weight
The TUE is used for transferring and attaching the PFU to the inner hub. Weight limit
which the TUE can carry is 3.0 tons and parts of HSC which is carried by the TUE should
weigh less than 3.0 tons. We once considered the possibility that we recreate the TUE inorder to be free from this tight constraint on HSC weight. However, we found that TUE
modiﬁcation is diﬃcult in both technical and cost aspects. Therefore, we have to meet
the HSC weight constraint of 3.0 tons. This does not include the ﬁlter exchanger which is
attached separately.

8.3. MECHANICAL MODIFICATION OF TELESCOPE

219

Figure 8.18: Telescope balancer positions and parameter deﬁnition in the balance condition.

Table 8.2: The current C/W positions (actual measurement) and top unit gravity center
positions (calculated). The red characters are actual measurements while blue characters
are calculated values.
No.
1
2
3
4
5
6

top unit
CAS inst.
CS-OPT M2
FOCAS
NS-OPT M2
MOIRCSIRM2
MOIRCS
IRM2+LLTMOIRCSPFU
MOIRCSFMOS
MOIRCS

top unit[kg]
2136.52405.02239.52577.02786.5?

balancerC/WbalancerC/WbalancerC/WbalancerC/WbalancerC/Wbalancer
C/W

A
420
138.8
202
-79.2
352
70.8
20
-261.2-487-768.2-366-647.2balancer weight
B
C
423
420
-177.5-180.5228
228-372.5-372.5352
340-248.5-260.53131-569.5-569.5-486-486-1086.5-1086.5-353-354-953.5-954.5

D
420
138.8
203
-78.2
339
57.8
32
-249.2-474-755.2-366-647.2

tel.
M0
-22.5

balancerMb-0.31

top unit
Mt
22.8

total
M
4.78E-06

top unit
GC
10656

-22.5

-3.43

25.9

-6.8E-06

10765

-22.5

-1.45

23.9

-1.9E-05

10675

-22.5

-6.27

28.7

-5.3E-06

11148

-22.5

-14.0536.5-1.7E-05

13101

-22.5

-12.17

34.6

0

?

220

CHAPTER 8. TELESCOPE INTERFACE

Constraint from the telescope balancer
We consider a telescope balance condition (equations below, Equation 8.1 to 8.8) about theelevation axis (Figure 8.18) in order to estimate a limit on the total weight of HSC.
Telescope balance conditions are
M1 + M2
M0
M1
M2
M3
M4
M5
W

=
=
=
=
=
=
=
:

−(M0 + M3 + M4 + M5 )W0 × L0 : current PFU
W1 × L1 : HSC in total
W2 × L2 : added to top ringW3 × L3 : C/W to balance with W2 /L2W4 × L4 : C/W to oﬀset C/W location of balancer
W5 × L5 : balancer C/Wweight, L : gravity center location

(8.1)
(8.2)
(8.3)
(8.4)
(8.5)
(8.6)
(8.7)

In these equations, M2 = (−1) × M3 . Therefore,
M1 = −(M0 + M4 + M5 )(8.8)

is satisﬁed.
There are four telescope balancers whose respective weights are 3.8 tons and adjustablestrokes are ±680 mm (see Figure 8.2). For all the current cases, of course, positions of thebalancers used are within the strokes of the balancers. Estimation error of these momentums
is 1 ton-m which corresponds to a balancer position oﬀset of ±66 mm. We here assume thatthe HSC barycentric position is the same as that of the current PFU. If we just add ﬁxed
counterweight (C/W) and do not modify the telescope balancer, stroke of the telescopebalancer used is ±614 mm. In this case, the weight limit for the HSC unit including theﬁlter exchange unit attached separately is 3.16 tons (see Table 8.3). If we add C/Ws andadd weights to the secondary mirror units (up to 3 tons), weight limit on the HSC unit is3.85 tons (see Table 8.4).
In these two cases, we have simulated frequency response characteristics of the telescopeas shown in Figure 8.19. For the HSC 3.16-ton case, there are only small diﬀerences between
the current PFU and HSC. On the other hand, we can see a slight diﬀerence at a frequencyof ∼ 60 Hz between the current PFU and HSC. However, the diﬀerence is as small as
the diﬀerence between the current secondary mirrors and the current PFU. Therefore, we
can now conclude that a comparable telescope tracking accuracy can be achieved for HSC
although we need more examinations.
For both cases, we do not need to modify the telescope balancer but do need to add ﬁxed
C/W, which enable us to balance the telescope with all the instruments. Details are shownin the document [9].

8.4

Schedule for telescope modiﬁcation

As described above, the telescope has to be modiﬁed to accommodate HSC onto the prime
focus to achieve the ﬁeld of view as required from the scientiﬁc goals as well as to satisfy theweight limit set by the TUE unit loader. In this section, we describe a possible schedule tomodify the telescope including down time expected in summers of 2010 and 2011. The goal8.4. SCHEDULE FOR TELESCOPE MODIFICATION

221

Table 8.3: Weight constraint on HSC when we add ﬁxed C/W without modifying the telescope balancer. In this case, HSC weight can be up to 3.16 tons. In ”total” column,
M = M0 + M5 + M1 + M4 . The balancer position columns indicate that the range of thebalancer positions used is ±614 mm. The last column, top unit GC, is measured with reference to the elevation axis in milimeter. We assume here that the HSC barycentric positionis the same as that of the current PFU.
No.
1
2
3
4
5
6
7

top unit
CAS inst.
CS-OPT M2
FOCAS
NS-OPT M2
MOIRCSIRM2
MOIRCS
IRM2+LLTMOIRCSPFU
MOIRCSFMOS
MOIRCS
HSC
MOIRCS

top unit[kg]
2136.52405.02239.52577.02786.5?
3160

balancerC/WbalancerC/WbalancerC/WbalancerC/WbalancerC/WbalancerC/Wbalancer
C/W

A
611
329.8
393
111.8
543
261.8
211
-70.2
-296
-577.2-175-456.2-614-895.2balancer weight
B
C
614
611
13.5
10.5
419
419
-181.5-181.5
543
531
-57.5
-69.5
222
222
-378.5-378.5-295-295-895.5-895.5-162-163-762.5-763.5-614-614-1214.5-1214.5

D
611
329.8
394
112.8
530
248.8
223
-58.2
-283
-564.2-175-456.2-614-895.2

tel.
M0
-22.5

balancerM5
2.60

top unit
M1
22.8

add C/W
M4
-2.9

total
M
1.1E-05

top unit
GC
10656

-22.5

-0.53

25.9

-2.9

-5.4E-07

10765

-22.5

1.46

23.9

-2.9

-1.3E-05

10675

-22.5

-3.36

28.7

-2.9

-9.6E-07

11148

-22.5

-11.1436.5-2.9-1.1E-05

13101

-22.5

-9.27

34.6

-2.9

-6.3E-06?

-22.5-16.0341.4-2.9

0.0E+00

13101

Table 8.4: Weight constraint on HSC when we add C/W to the secondary mirrors without
modifying the telescope balancer. In this case, HSC weight can be up to 3.85 tons. In ”total”
column, M = M0 + Mb + Mt + Ma . The balancer position columns indicate that the rangeof the balancer positions used is ±614 mm. The last column, top unit GC, is measured
with reference to the elevation axis in milimeter. We assume here that the HSC barycentric
position is the same as that of the current PFU.
No.
1
2
3
4
5
6
7

top unit
CAS inst.
CS-OPT M2
FOCAS
NS-OPT M2
MOIRCSIRM2
MOIRCS
IRM2+LLTMOIRCSPFU
MOIRCSFMOS
MOIRCS
HSC
MOIRCS

top unit
[kg]
3000
3000
3000
3337.52786.5?
3851.5balancerC/WbalancerC/WbalancerC/WbalancerC/WbalancerC/WbalancerC/Wbalancer
C/W

A
599.3
318.1
585.6
284.4
602.7
321.5
247.1
-34.1
300.2
19.0
421.2
140.0
-614
-895.2balancer weight
B
C
602.3
599.3
1.8
-1.2
591.6
591.6
-8.9
-8.9
602.7
590.7
2.2
-9.8
258.1
258.1
-342.4-342.4301.2
301.2-299.3-299.3434.2
433.2-166.3-167.3-614-614-1214.5-1214.5

D
599.3
318.1
566.6
285.4
589.7
308.5
259.1
-22.1
313.2
32.0
421.2
140.0
-614
-895.2

tel.
M0
-22.5

balancerMb
2.42

top unit
Mt
32.0

add C/W
Ma
-12.0

total
M
-3.8E-02

top unit
GC
10656

-22.5

2.10

32.3

-12.0

-3.4E-02

10765

-22.5

2.37

32.0

-12.0

-3.7E-02

10675

-22.5

-2.82

37.2

-12.0

-3.5E-02

11148

-22.5

-2.08

36.5

-12.0

-1.2E-06

13101

-22.5

-0.20

34.6

-12.0

-2.0E-07?

-22.5-16.03

50.5

-12.0

0.0E+00

13101

222

CHAPTER 8. TELESCOPE INTERFACE

Figure 8.19: Frequency response characteristics of the telescope elevation axis. We show
six cases of the current CS-Opt (red), the current NS-Opt (light purple), the current IRM2
(green), the current PFU (yellow), the HSC 3.16 ton case (black), and the HSC 3.85 toncase (purple).

8.4. SCHEDULE FOR TELESCOPE MODIFICATION

223

Table 8.5: Additional tasks expected while modiﬁcation of the telescope for HSC
WBS number
3.0
4.7.3

SummaryEﬀort
Prepare for re-aluminization
5.9 man-months
Maintain primary mirror actuators 28 man-daysof this section is to establish a reasonable schedule to modify the telescope and commission
HSC
1. within reasonable workload on engineers and technicians,
2. meeting the speciﬁcations,
3. within the expected time frame and cost, and4. within acceptable impact to scientiﬁc observations using the telescope.
As shown below, the baseline schedule expects down time of at least 3.0 ± 0.5 months
in summer 2010 and 2.5 ± 0.5 months in summer 2011. The down time in 2010 might belonger if workload to implement collision prevention measures (Tables 8.8 and 8.9) is high.
After the down time in 2011, 1.5 ± 0.5 months will be dedicated to commission HSC.

8.4.1

Tasks

Tasks needed to accommodate HSC meeting the speciﬁcations are listed in a Gantt chart[10].
The chart covers tasks to modify the telescope and existing top units mechanically, add cables
and controllers, test the modiﬁed and added features of the telescope, and commission HSC.
In this study, work breakdown structure (WBS) numbers are assigned following the Gantt
chart[10].
In addition to the tasks needed to accommodate HSC, tasks for re-aluminization of the
primary mirror and regular maintenance work while re-aluminization have to be taken into
account. To modify the inner hub, the telescope primary mirror needs to be removed toerect a scaﬀold for the machine tool. To minimize impact to scientiﬁc observations, realuminization of the telescope primary mirror is scheduled in summer 2010, to be done when
the telescope is being modiﬁed for HSC. Table 8.5 lists such additional tasks that are notincluded in the Gantt chart[10]. The preparatory work for re-aluminization (WBS number
3.0) is also included to account for its workload while preparing for telescope modiﬁcation
tasks while re-aluminization. The task to modify LLT (WBS number 2.5.4) has been removed
because this task can be and should be done in 2009.
Workload of each tasks to be done by the contractor are taken as is from the Gantt
chart[10]. As for telescope modiﬁcation tasks to be done by the staﬀ members of the observatory, we reevaluated workloads to reﬂect our own experiences. Workload for tasks related
to re-aluminization of the telescope primary mirror are estimated from records of the previous
re-aluminization in summer 2006.
Design and test of the collision prevention measures listed in Tables 8.9 are not taken
into account in this study. However, we foresee the following steps to ensure safety of thelens tube. In 2009, while tests going on to evaluate collision prevention measures, interface

224

CHAPTER 8. TELESCOPE INTERFACE

between sensors, controls including feed backs and interlocks, as well as responsible parties
in implementation are to be clariﬁed. In summer 2010, the successful measures to prevent
collision should be installed. The collision prevention measures are to be adjusted whileregular top unit exchanges with the existing top units to establish safe sequences for the lens
tube before HSC comes onto the telescope.

8.4.2

Resources

The modiﬁcation of the telescope and commissioning of HSC will be a collaborative work of
a number of contractors and a number of branches of National Astronomical Observatory ofJapan (NAOJ) including Advanced Technology Center and Subaru Telescope.
At Subaru Telescope, we expect to have six engineers from telescope engineering divisionand seven technicians as day crews. To keep the telescope operational, 3.8 engineers andtwo technicians are constantly maintaining the telescope while others taking care of tasks
such as upgrades of the telescope and supporting instruments that are being commissioned.
When re-aluminizing the primary mirror, other staﬀ members of the observatory, including
scientists and administrative oﬃcers, help complete many tasks for re-aluminization andtelescope maintenance to be done. In this study, eight staﬀ members, possibly including
those from other branches of NAOJ, are assumed to help summit work in July 2010 whenre-aluminization is scheduled.
It should be noted that the summit of Mauna Kea, where the telescope resides, is 4,200
meters high. The low partial pressure of oxygen (60% of the sea level) and freezing temperature inside the dome exhaust workers rapidly. For this study, engineers and technicians are
assumed to work four days a week at the summit to limit workload within reasonable limit.

8.4.3

Estimated scheduleWith taking into account the tasks and the resources stated above, the tasks are scheduledusing a project management tool3 which schedules tasks to comfort resource limits. As
described in section 8.4.3.1, accuracy of scheduling is estimated to be a half month on
duration of summary tasks. Table 8.6 shows the resulting schedule while details are shownin Figures 8.20 – 8.23.
In summer 2010, telescope down time of three months from July 1 through September
28 is expected. This might be longer in case if workload to implement collision prevention
measures is high. In summer 2011, the telescope will not be on the sky for a period of 2.5
months from July 1 through September 12 which is followed by 1.5 months of engineeringobservations from September 13 through October 26 to commission HSC.
With many engineers and technicians from diﬀerent branches and contractors working on
many interrelated tasks simultaneously, overall management of the resources and tasks isimportant to eﬀectively modify the telescope and commission HSC.
8.4.3.1

Accuracy

Accuracy of the schedule discussed here is estimated to be about a half month in duration.
3

TaskJuggler http://www.taskjuggler.org
Start date of WBS 4: Down time (2010) excludes WBS 4.6.1.1: Carry-in of material for scaﬀold, which
begins before stopping observations.
48.4. SCHEDULE FOR TELESCOPE MODIFICATION

225

Gantt Chart toward commisioning of HSC (2010)
Name
Maintain telescope
2 Tasks before downtime(2010)
2.1 頭頂外環まわり新設ケーブル布設(el5度)
2.2 頭頂外環まわり新設ケーブル布設(天頂)2.3 ELカウンターウェイト用センターセクション加工
2.4 望遠鏡上電気系端子架新設用タップ加工
2.5 副鏡関連改修
2.5.1 副鏡ユニットの養生、タップ加工、ウェイト追加(2台分)
2.5.2 副鏡ユニットのジャンパーケーブル追加(2台分)
2.5.3 副鏡ユニットのジャンパーケーブル経由試験(2台分)
3 Realuminization
3.0 Preparation
3.0.1 Ventilation
3.0.1.2 Installation of fans
3.0.3 Meeting
3.0.4 Check for fans inside Cs hole
3.0.5 Tool inventory
3.0.6 Move M1 trolley to move boom lift
3.0.7 Lower small boom lift
3.0.9 PLumb mirror wash area
3.0.10 Maintain pure water facility
3.0.11 Check and maintain 80t crane
3.0.12 Clean up 80t crane
3.0.13 Maintain M1 trolley
3.0.14 Clean ESB
3.0.15 Vacuum clean inside chamber
3.0.16 カセグレンプラグ動作試験
3.0.17 Install and test filaments
3.0.18 Setup monitor PC to chamber
3.0.19 Replace deposition controller
3.0.20 Drive test lower chamber
3.0.21 Test glow discharge
3.0.21.1 Test
3.0.21.2 Study
3.0.21.3 Study glow rod
3.0.21.4 Study argon flow
3.0.21.5 Confirm repair
3.0.22 Test heating all filaments
3.0.23 Pump up waste water from pit under ESB
3.0.24 Install thermometer in ESB
3.0.25 ベンディング流量調整用ジグ製作
3.0.26 Install quick connect for PMWF manifold
3.0.27 Fix rough vacuum pump
3.0.28 Safety issue
3.0.29 Check batteries on cassegrai trolley
3.0.30 Maintain M1 washing facility
3.0.30.1 Maintenance3.0.30.2 Remove drip pan for repair
3.0.30.3 Repiar drip pan
3.0.30.4 Clean up
3.0.32 Check N2 manifold
3.0.33 Work on ESB trolley
3.0.34 Clean up cell trolley
3.0.35 Maintain mirror hatches
3.0.35.1 Clean
3.0.35.2 Grease drive shafts
3.0.36 Modify portable step for vacuum chamber
3.0.37 Maintain M1 chamber
3.0.38 Tasks just before removal of M1
3.0.38.1 Measure M1 reflectivity
3.0.38.2 Move TWS
3.0.38.3 Wash sample glass plates
3.0.38.4 Disassemble lower cover3.0.38.5 Prepare HCl into tank
3.0.38.6 Prepare HCl swabs
3.1 主鏡取外し
3.1.1 Detach mirror
3.1.2 Work on CIAX
3.1.3 Prepare for washing
3.2 主鏡再蒸着
Container Task
Milestone

Start
End
2010$04$01 2010$06$29
2010$05$14 2010$06$30
2010$05$18 2010$05$25
2010$06$04 2010$06$09
2010$05$25 2010$06$03
2010$06$16 2010$06$18
2010$05$14 2010$06$08
2010$05$14 2010$05$26
2010$05$26 2010$06$01
2010$06$01 2010$06$08
2010$04$16 2010$08$17
2010$04$16 2010$06$30
2010$06$10 2010$06$15
2010$06$10 2010$06$10
2010$06$30 2010$06$30
2010$06$04 2010$06$07
2010$06$08 2010$06$10
2010$06$15 2010$06$15
2010$06$15 2010$06$15
2010$06$09 2010$06$15
2010$06$09 2010$06$15
2010$06$08 2010$06$18
2010$06$22 2010$06$24
2010$06$03 2010$06$10
2010$06$15 2010$06$18
2010$04$16 2010$04$16
2010$06$24 2010$06$25
2010$04$20 2010$04$20
2010$06$25 2010$06$25
2010$04$16 2010$04$16
2010$04$20 2010$04$20
2010$04$21 2010$06$04
2010$04$21 2010$04$21
2010$04$21 2010$05$07
2010$05$07 2010$05$19
2010$05$19 2010$06$01
2010$06$01 2010$06$04
2010$06$04 2010$06$10
2010$06$29 2010$06$29
2010$06$01 2010$06$02
2010$06$10 2010$06$18
2010$06$24 2010$06$29
2010$06$07 2010$06$23
2010$06$22 2010$06$24
2010$06$28 2010$06$29
2010$06$03 2010$06$24
2010$06$03 2010$06$07
2010$06$22 2010$06$23
2010$06$23 2010$06$23
2010$06$23 2010$06$24
2010$06$24 2010$06$25
2010$06$24 2010$06$25
2010$06$25 2010$06$25
2010$06$17 2010$06$22
2010$06$17 2010$06$17
2010$06$17 2010$06$22
2010$06$25 2010$06$29
2010$06$29 2010$06$30
2010$06$29 2010$06$30
2010$06$30 2010$06$30
2010$06$30 2010$06$30
2010$06$29 2010$06$30
2010$06$29 2010$06$30
2010$06$30 2010$06$30
2010$06$30 2010$06$30
2010$07$01 2010$07$08
2010$07$01 2010$07$07
2010$07$07 2010$07$08
2010$07$01 2010$07$07
2010$07$07 2010$07$20
Planned Task
In$progress Task

Apr

May

Completed Task
Task Work Load

Jun

2010
Jul

Aug

Sep

Oct

Other Work Load
Free Load

Page 1 of 2

2009$02$26
Generated on 2009$02$26 with TaskJuggler 2.4.1

Figure 8.20: Baseline schedule for 2010

226

CHAPTER 8. TELESCOPE INTERFACE

Gantt Chart toward commisioning of HSC (2010)
Name
Start
End
3.2.1 Inspect scratch
2010$07$07 2010$07$09
3.2.2 Wash
2010$07$09 2010$07$12
3.2.3 Inspect scratch
2010$07$13 2010$07$13
3.2.4 Repair scratch
2010$07$14 2010$07$14
3.2.5 Wash
2010$07$14 2010$07$15
3.2.6 Seal boggie
2010$07$16 2010$07$20
3.2.7 Evapolate
2010$07$15 2010$07$16
3.3 主鏡取付け
2010$08$13 2010$08$17
3.3.1 Attach mirror
2010$08$13 2010$08$16
3.3.2 Attach Cs instrument
2010$08$16 2010$08$17
4 Down time (2010)2010$06$29 2010$09$29
4.1 トップユニット交換(最終観測時のトップユニット撤去
2010$07$01 2010$07$01
4.2 副鏡ユニットの養生、タップ加工、ウェイト追加(1台分)
2010$07$01 2010$08$12
4.3 副鏡ユニットのジャンパーケーブル追加(1台分)
2010$08$12 2010$08$18
4.4 副鏡ユニットのジャンパーケーブル経由試験
2010$08$18 2010$08$18
4.5 TUE改造
2010$07$01 2010$08$02
4.5.1 シーケンサ交換、改修版SWでの単体試験
2010$07$01 2010$08$02
4.6 頭頂内環改修
2010$06$29 2010$08$16
4.6.1 足場設置
2010$06$29 2010$07$19
4.6.1.1 足場材搬入
2010$06$29 2010$07$01
4.6.1.2 第3鏡養生
2010$07$08 2010$07$19
4.6.1.3 足場設置
2010$07$01 2010$07$08
4.6.1.3.1 M3までの足場
2010$07$01 2010$07$07
4.6.1.3.1 M3以上の足場
2010$07$01 2010$07$08
4.6.2 機械系工事
2010$07$06 2010$07$27
4.6.2.1 準備
2010$07$06 2010$07$08
4.6.2.2 頭頂内環内のフランジ切除
2010$07$08 2010$07$19
4.6.2.3 頭頂内環内の新クリップ機構タップ加工
2010$07$19 2010$07$26
4.6.2.4 頭頂内環内の新クリップ機構用ダミーウェイト設置
2010$07$26 2010$07$27
4.6.3 電気系工事
2010$07$27 2010$08$16
4.6.3.1,2 副鏡自動脱着コネクタ部の改造・コネクタ設置・配線検査 2010$07$27 2010$08$10
4.6.3.3 Tip/tiltケーブルの天から地への移設 (el5)2010$08$13 2010$08$16
4.6.4 足場解体
2010$08$10 2010$08$13
4.7 その他望遠鏡改修
2010$07$01 2010$08$12
4.7.1 機械系工事
2010$07$27 2010$07$30
4.7.1.1 望遠鏡上電気系端子架の設置
2010$07$27 2010$07$27
4.7.1.2 センターセクションのカウンターウェイト設置
2010$07$27 2010$07$30
4.7.2 電気系工事
2010$08$10 2010$08$12
4.7.2.1 電気系端子架関連ケーブルの端末処理
2010$08$10 2010$08$12
4.7.3 主鏡アクチュエータの保守
2010$07$01 2010$07$13
4.7.3.1 アクチュエータの交換
2010$07$01 2010$07$02
4.7.3.2 アクチュエータモーターの交換
2010$07$01 2010$07$13
4.8 片付け2010$08$18 2010$08$23
4.9 IRM2(CsあるいはNsIR ImRなし)での望遠鏡復旧
2010$08$16 2010$09$29
4.9.1 全系通し配線検査
2010$08$16 2010$08$19
4.9.2 トップユニット交換 (→IRM2)
2010$08$19 2010$09$02
4.9.2.1 副鏡実装、手動ケーブル接続、落下防止確認
2010$08$19 2010$09$02
4.9.3 副鏡制御再現確認
2010$09$02 2010$09$10
4.9.3.1 安全確認
2010$09$02 2010$09$07
4.9.3.2 機能、性能の再現確認
2010$09$07 2010$09$10
4.9.4 望遠鏡バランス調整
2010$09$10 2010$09$15
4.9.5 望遠鏡追尾精度再現確認、応答特性データ取得
2010$09$15 2010$09$21
4.9.6 PA、MA
2010$09$21 2010$09$29

Container Task
Milestone

Planned Task
In$progress Task

Apr

May

Completed Task
Task Work Load

Jun

2010
Jul

Aug

Sep

Oct

Other Work Load
Free Load

Page 2 of 2

2009$02$26
Generated on 2009$02$26 with TaskJuggler 2.4.1

Figure 8.21: Continued baseline schedule for 2010

8.4. SCHEDULE FOR TELESCOPE MODIFICATION

227

Gantt Chart toward commisioning of HSC (2011)
Name
Start
End
Maintain telescope
2011$04$01 2011$05$31
5 Tasks before down time (2011)2011$04$11 2011$06$30
5.1 観測停止前の工事
2011$05$31 2011$06$30
5.1.1 準備、開梱(製品、工具類、養生材他)
2011$06$16 2011$06$17
5.1.2 望遠鏡上作業
2011$06$16 2011$06$30
5.1.2.1 赤外反射防止ミラー用タップ加工
2011$06$28 2011$06$30
5.1.2.2 ドームフラット光源移設
2011$06$27 2011$06$30
5.1.2.3 新クリップ機構制御部設置(TUE ULトロリー上)
2011$06$16 2011$06$17
5.1.2.4 新クリップ機構配線(TUE ULトロリー上)
2011$06$17 2011$06$22
5.1.2.5 新クリップ機構制御部の事前動作確認(TUE ULトロリー上) 2011$06$22 2011$06$22
5.1.2.6 頭頂内環外の新クリップ機構取付
2011$06$22 2011$06$23
5.1.2.7 赤外反射防止ミラー取付
2011$06$22 2011$06$22
5.1.2.8 頭頂内環外の新クリップ機構配線
2011$06$24 2011$06$28
5.1.2.9 新クリップ機構の事前クリップ動作確認
2011$06$30 2011$06$30
5.1.3 制御棟2回作業
2011$05$31 2011$06$30
5.1.3.1 機械系工事
2011$06$29 2011$06$30
5.1.4.1.1 制御棟2階新設筐体据付
2011$06$29 2011$06$30
5.1.3.2 電気系工事
2011$05$31 2011$06$30
5.1.3.2.1 新設筐体用電源ケーブルの布設・末端処理
2011$05$31 2011$06$23
5.1.3.2.2 新設ケーブル(一部)の布設・末端処理
2011$06$01 2011$06$13
5.1.3.2.3 片づけ
2011$06$29 2011$06$30
5.2 HSCユニットの現地組立
2011$04$11 2011$04$28
5.2.1 準備・開梱(HSCユニット、ジグ、試験ケーブル)
2011$04$11 2011$04$13
5.2.2 補正光学系組込、ケーブル・配管接続
2011$04$13 2011$04$15
5.2.3 USのHSC設置位置にかさ上げアダプタ取付
2011$04$15 2011$04$18
5.2.4 TUE US上への移動・取付
2011$04$18 2011$04$19
5.2.5 観測装置・フィルター・シャッタ組込・ケーブル・配管接続
2011$04$19 2011$04$28
6 Down time (2011)2011$07$01 2011$07$27
6.1 準備2011$07$01 2011$07$05
6.2 望遠鏡上作業
2011$07$01 2011$07$13
6.2.1 トップユニット交換(最終観測時のトップユニット撤去)
2011$07$01 2011$07$01
6.2.2 TUE改造
2011$07$01 2011$07$08
6.2.2.1 US、UTのHSC搬送動作試験
2011$07$01 2011$07$06
6.2.2.2 ULのHSC搬送動作試験
2011$07$06 2011$07$08
6.2.3 機械系工事
2011$07$08 2011$07$13
6.2.3.1 HSC実装・機械系フィットチェック(クリップ動作確認含む) 2011$07$08 2011$07$12
6.2.3.2 HSC取外し
2011$07$12 2011$07$13
6.3 制御棟2階作業
2011$07$05 2011$07$27
6.3.1 電気系工事
2011$07$05 2011$07$27
6.3.1.1 準備
2011$07$06 2011$07$06
6.3.1.2 制御棟内切替用ケーブルの布設・末端処理
2011$07$12 2011$07$18
6.3.1.3 新設制御機器と既設制御機器間のケーブル布設・末端処理 2011$07$07 2011$07$11
6.3.1.4 布設部分の配線検査
2011$07$19 2011$07$25
6.3.1.5 片付け2011$07$26 2011$07$27
7 Commissioning of HSC (not on telescope)
2011$04$28 2011$07$20
7.1 天文台担当計算機$HSCユニット間組合せ試験
2011$04$28 2011$05$31
7.1.1 TUE階試験(TUE階用試験ケーブル要)
2011$04$28 2011$05$31
7.2 制御計算機試験
2011$05$13 2011$07$20
7.2.1 計算機単体試験
2011$05$13 2011$06$15
7.2.1.1 S/Wインストール、起動確認、計算機間通信確認
2011$05$13 2011$05$18
7.2.1.2 計算機間機能確認試験
2011$06$14 2011$06$15
7.2.2 ローカル制御装置との組合せ試験
2011$06$15 2011$06$28
7.2.2.1 通信試験
2011$06$27 2011$06$28
7.2.3 天文台担当品との組合せ
2011$06$29 2011$07$12
7.2.3.1 通信試験
2011$06$29 2011$07$06
7.2.3.2 リモート監視機能確認試験
2011$07$07 2011$07$12
7.2.4 HSCユニット@TUE$US上
2011$05$19 2011$05$31
7.2.4.1 リモート監視機能確認試験
2011$05$24 2011$05$25
7.2.4.2 リモート制御機能確認試験
2011$05$26 2011$05$31
7.2.4.3 計算機インタロック機能確認
2011$05$19 2011$05$20
7.3 ローカル制御装置$HSCユニット組合せ試験(TUE$US上)
2011$05$31 2011$07$20
7.3.1 全系通し配線検査(新設切替架経由の配線含む)
2011$06$01 2011$06$03
7.3.2 HSCユニット組合せ(TUE$US上、EL駆動中は不可)
2011$06$03 2011$07$20
7.3.2.1 インストルメントロテータ制御
2011$06$03 2011$06$13
7.3.2.1.1 機能確認(リミット動作、駆動範囲)
2011$06$06 2011$06$13
7.3.2.2 ヘキサポッド制御
2011$06$03 2011$06$23
7.3.2.2.1 機能確認(リミット動作、駆動範囲)
2011$06$16 2011$06$23
7.3.2.3 ADC制御
2011$06$03 2011$07$20
7.3.2.3.1 機能確認(リミット動作、駆動範囲)
2011$07$13 2011$07$20
8 Commissioning of HSC (on telescope in day time)
2011$07$20 2011$09$13
Container Task
Milestone

Planned Task
In$progress Task

Apr

May

Completed Task
Task Work Load

Jun

2011
Jul

Aug

Sep

Oct

Other Work Load
Free Load

Page 1 of 2

2009$02$26
Generated on 2009$02$26 with TaskJuggler 2.4.1

Figure 8.22: Baseline schedule for 2011

228

CHAPTER 8. TELESCOPE INTERFACE

Gantt Chart toward commisioning of HSC (2011)
Name
8.1 HSCユニット実装、バランス調整
8.1.1 副鏡→HSC実装、手動ケーブル接続、落下防止確認8.1.2 望遠鏡バランス調整
8.2 ローカル制御装置試験$HSCユニット組合せ試験(望遠鏡上)
8.2.1 HSCユニット組合せ(望遠鏡上)
8.2.1.1 インストルメントロテータ制御
8.2.1.1.1 機能確認(リミット動作、駆動範囲)
8.2.1.2 ヘキサポッド制御
8.2.1.2.1 機能確認(リミット動作、駆動範囲)
8.2.1.3 ADC制御
8.2.1.3.1 機能確認(リミット動作、駆動範囲)
8.3 天文台担当計算機$HSCユニット間組合せ試験
8.3.1 望遠鏡上試験
8.4 制御計算機試験
8.4.1 ローカル制御装置との組合せ試験
8.4.1.1 HSCユニット@望遠鏡上
8.4.1.1.1 EL連動リモート制御、監視機能確認試験
8.4.2 天文台担当品との組合せ
8.4.2.1 リモート制御機能確認試験
8.4.2.1.1 AGリモート制御試験
8.4.2.1.2 SHリモート制御試験
8.4.2.1.3 SH専用AGリモート制御試験
8.4.3 シーケンス確認
8.4.3.1 昼間のシーケンスの確認(望遠鏡指向、追尾)
8.4.3.2 昼間のPAシーケンスの確認
8.4.3.2 昼間のMAシーケンスの確認
9 Commissioning of HSC (on sky)
9.1 現地総合試験
9.1.1 観測天体指向、オープンループ追尾(EL連動含む)
9.1.2 AGによるガイド星捕捉、追尾誤差の座標確認
9.1.3 AGクローズドループ追尾
9.1.4 SHによる波面検出、波面誤差の座標確認
9.2 総合光学調整
9.2.1 HSCのPA、MA蚶補正取得
9.2.2 HSCのPA、MA性能追い込み
9.2.3 望遠鏡クローズドループ追尾精度確認
9.2.4 AZ/EL駆動のフィルター調整

Container Task
Milestone

Start
End
2011$07$20 2011$07$27
2011$07$21 2011$07$26
2011$07$27 2011$07$27
2011$07$27 2011$08$11
2011$07$27 2011$08$11
2011$07$27 2011$08$08
2011$08$03 2011$08$08
2011$07$27 2011$08$11
2011$08$09 2011$08$11
2011$07$27 2011$08$02
2011$07$28 2011$08$02
2011$07$27 2011$08$02
2011$07$28 2011$08$02
2011$08$11 2011$09$13
2011$08$11 2011$08$16
2011$08$11 2011$08$16
2011$08$15 2011$08$16
2011$08$16 2011$09$01
2011$08$16 2011$09$01
2011$08$22 2011$08$25
2011$08$29 2011$09$01
2011$08$17 2011$08$18
2011$09$01 2011$09$13
2011$09$06 2011$09$07
2011$09$12 2011$09$13
2011$09$08 2011$09$09
2011$09$13 2011$10$26
2011$09$13 2011$10$03
2011$09$14 2011$09$15
2011$09$16 2011$09$20
2011$09$21 2011$09$26
2011$09$26 2011$10$03
2011$10$03 2011$10$26
2011$10$03 2011$10$10
2011$10$03 2011$10$06
2011$10$03 2011$10$11
2011$10$11 2011$10$26

Planned Task
In$progress Task

Apr

May

Completed Task
Task Work Load

Jun

2011
Jul

Aug

Sep

Oct

Other Work Load
Free Load

Page 2 of 2

2009$02$26
Generated on 2009$02$26 with TaskJuggler 2.4.1

Figure 8.23: Continued baseline schedule for 2011

8.4. SCHEDULE FOR TELESCOPE MODIFICATION

229

Table 8.6: Durations of tasks on baseline schedule
WBS number
2
3
4
5
6
7
8
9

Summary
Before down time (2010)
Re-aluminization (including preparation)
Down time (2010)4
Before down time (2011)
Down time (2011)
Commissioning (not on telescope)
Commissioning (on telescope, daytime)
Commissioning (on sky)

Start
End
2010-5-14 2010-7-1
2010-4-16 2010-8-17
2010-7-1 2010-9-29
2011-4-11 2011-7-1
2011-7-1 2011-7-27
2011-4-28 2011-7-20
2011-7-20 2011-9-13
2011-9-13 2011-10-26

Table 8.7: Diﬀerence of schedules due to project management tools

WBS number
2
3
4
5
6
7
8
9

Duration (days) Diﬀerence
Summary[10] This study(days)
Before down time (2010)
34
16
18
Re-aluminization
55
41
14
Down time (2010)
78
66
12
Before down time (2011)
33
42
9
Down time (2011)20
23
3
Commissioning (not on telescope)82
72
10Commissioning (on telescope, daytime)
42
84
42
Commissioning (on sky)
45
37
8
Average diﬀerence
14.5This study is done with a project management tool TaskJuggler. To test agreement with
the Gantt chart[10], we scheduled tasks following the chart as close as possible. Some of
the resulting durations of summary tasks are much diﬀerent. This seems to be from using
a diﬀerent project management tool and possibly from hidden dependencies and constraints
between tasks in the chart. Table 8.7 shows the diﬀerence. The average diﬀerence in durationof the summary tasks is a half month.

8.4.4

Impacts on operation of observatory

A number of impacts on operation of the observatory due to assigning all the available
resources to modiﬁcation of the telescope and commissioning of HSC should be addressed
especially in spring – summer of 2010 and in spring 2011 when workload on engineers andtechnicians are high.
As shown below, these impacts can be minimized by planning ahead or delaying tasks not
related to commissioning of HSC.

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8.4.4.1

CHAPTER 8. TELESCOPE INTERFACE
Delay in telescope upgrades

One of the primary tasks for the engineers at the observatory has been to propose, plan,
design, and implement new features of the telescope, in collaboration with a contractor if
necessary. Telescope upgrades should be planned ahead.
8.4.4.2

Unable to support to instruments

From time to time, the engineers at the observatory have been supporting instrument developers, providing with interface information, advising from their experiences, and working at
the telescope. Similarly to the telescope upgrades, necessary tasks should be planned aheador otherwise delayed.
8.4.4.3

Unable to help top unit exchangeFirst, the laser launch telescope (LLT) still requires two engineers to be mounted or detached.
The workload should be integrated in the schedule to allow the engineers to work on LLTwhen needed.
Second, top unit exchange will require more technicians to be completed after summer
2010. Removal of the auto-connectors which are being used to electrically connect thesecondary mirrors to the telescope control system requires technicians to manually connect
and disconnect electrical connectors. Additional workload should be incorporated into theschedule for modiﬁcation of the telescope and commissioning of HSC.
The schedule shown in this study does not include workload for LLT and additional workload after removal of auto-connectors. These workloads might slightly extend duration ofsome of the tasks however impact of these workloads to modiﬁcation of the telescope andcommissioning of HSC can be minimized by carefully scheduling observations and top unit
exchanges.

8.5

Safety Issues on Attachment of HSC to Telescope

8.5.1

Concern on regular operations

The wide ﬁeld corrector (WFC) barrel is made of Cordierite which is a type of ﬁne ceramics.
The ﬁne ceramics is high-speciﬁc stiﬀness material, however the brittleness is generally higher
than that of metal. Over ten years operation of the Subaru telescope, there were a few
collision cases between the top unit and the telescope during transferring the top unit. The
concern is that kind of collision leads to a fracture of the WFC barrel and a secondary
accident especially on the primary mirror.

8.5.2

Top unit exchange sequence

The top units including HSC are transferred from a storage to the inner hub of the telescopeby the top unit exchanger (TUE). The top unit exchange work has been performed asshown in Figure 8.24. In the ﬁgure, a secondary mirror is grabbed by the unit loader andapproaching to the inner hub. The expected sequence of attaching HSC is summarized in
Figure 8.25. HSC will be stored at the top unit storage ﬂoor as well as the other top units.
First, HSC is taken from the storage ﬂoor. Next, the unit loader, the robot arm on the girder

8.5. SAFETY ISSUES ON ATTACHMENT OF HSC TO TELESCOPE

Girder Bridge

231

Top Ring

Top Unit

Inner Hub

Storage Floor
Unit Loader

Secondary mirror

Figure 8.24: One scene during TUE work. A secondary mirror is approaching to the innerhub.
bridge, picks up the unit. The unit loader retracts its arm to uplift the unit and moves toright above the inner hub. The operators make precise adjustments manually to align theunit to the inner hub, then lower the unit carefully and clamp to the telescope.
As Figure 8.26 shows, the minimum clearance in radius will be 15 mm at the ADC ﬂange,
next one will be 22 mm at the G1 ﬂange[11] In addition, there remains uncertainty of 5
mm of the unit loader vibration caused by operator activity and gust wind[12]. Because
the clearance between the WFC barrel and the inner hub is not so large, there are twoconcerns while operating TUE due to ﬂexure of the unit loader arm. One concern is tilt of
the unit[13]. The tilt can be more than 0.4 degrees by weight comparison with existing top
units. If this tilt will not be corrected, the minimum clearance will be less than describedabove. Second concern is horizontal drifts of the unit loader while inserting HSC into theinner hub. Because the horizontal drift needs to be compensated by manual operation, whichis considered to be the most dangerous operation in the top unit exchange sequence because
there is the risk of human error.

8.5.3

Contact stress analysis

The contact stress analysis is made by estimating the contact force at collision. To translate
force to stress, a contact area must be required. By assuming Herzian contact , the contact
area can be obtained by
16 r1 r2 1 − ν12 1 − ν22
(
+
)F h,
A =
π r1 − r2 E1
E2
2

(8.9)

where A is the contact area, r1 is the inner radius of the inner hub, r2 is the radius of theWFC barrel, ν is Poisson ratio, E is Young’s modulus, F is applied force, and h is theheight of the contact area. The suﬃx 1 and 2 denote the inner hub and the WFC barrel,
respectively. Then, we can compare the force estimated by the case study with allowable

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CHAPTER 8. TELESCOPE INTERFACE

Unit Loader

HSC

Inner Hub

1. Take out HSC

2. Hand over tofrom storage

Unit Loader

3. Align and insert HSC
into inner hub

Figure 8.25: Expected sequence of attaching HSC.

15 mm
in radius at
ADC flange

22 mm
in radius at
G1 flange

Figure 8.26: Clearance on the middle way of installation.

8.5. SAFETY ISSUES ON ATTACHMENT OF HSC TO TELESCOPE

233

contact stress. Also, we focus on the tensile stress in this analysis because the allowabletensile stress of ceramics is smaller than the allowable compressive stress.
Next, we need to deﬁne the safety factor. The probability of fracture is considered tofollow the Weibull distribution and given by
Pf = 1 − exp[−(

1 m
) ],
Sf

(8.10)

where Pf is probability of fracture, Sf is safety factor and m is Weibull modulus. Weibull
modulus corresponds to the variability in measured strength and reﬂects the distribution of
ﬂaws in the material. To increase the reliability of the material, we add the process called
Hot Isostatic Pressing (HIP) which reduces the root of fracture such as voids or ﬂaws. By
the HIP process, the Weibull modulus of Cordierite can be increased from 9.6 to 14.4[14].
Kyocera, the vendor of Cordierite, has limited the probability of 10−6 by the speciﬁcation
for the maximum probability of fracture, therefore the safety factor must be larger than 2.6.
In addition, the actual allowable stress depends on the volume where the stress is applied.
This is called the size eﬀect and described as
VE1 1/m) ,(8.11)
VE2
where µ is allowable stress, VE is volume and m is Weibull modulus. The suﬃx 1 and 2
denote the test piece and the actual piece, respectively. The following studies take this size
eﬀect into account.
µ2 = µ1 (8.5.3.1

Collision case study (1)One of the collision case is the vibration of the unit loader. It is assumed that the unit
loader makes simple harmonic motion with an inertia mass of 9,768 kg which is sum of theunit loader and the HSC unit, and vibration speed of 2.5 mm/sec. The vibration speed has
not been measured directly, and is estimated by operators based on their experience. The
calculated force by the collision is 3,430 N and the estimated safety factor is 2.9[15]. Although
2.9 is above the required value (2.6), the safety factor can be less than the estimation in thereal cases because this study does not take an oﬀset collision into account.
8.5.3.2

Collision case study (2)Another collision case is an erroneous operation. If the erroneous operation happens while
the WFC barrel is passing through the inner hub, the unit loader pushes the WFC barrel
to the inner hub by the static thrust. The rated thrust of the unit loader is 11,040 N andmaximum thrust will be 38,000 N[16]. Another report shows that the safety factor will be
below 2.3 if the stress above 10,300 N is applied to the G1 ﬂange[17]. Therefore, we have toprevent the unit loader from pushing the WFC barrel and limit the thrust of the unit loader.

8.5.4

Preventive and post-fracture measures

At the present moment, we have decided to make several mandatory measures as shown in
Table 8.8. The ﬁrst measure is adding wedges at grip points of the HSC unit. As mentioned
above, the tilt of the unit decreases eﬀective clearance in the inner hub. The wedges will
be inserted between the HSC unit and the unit loader so that the WFC barrel will be

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CHAPTER 8. TELESCOPE INTERFACE

Table 8.8: Mandatory preventive measures

Add wedges
Shock absorbing pads
Crash pad
TUE thrust limit
Wire cage

Keep the unit parallel to the inner hub
Reduce contact stress
Prevent propagation of contact stress
Limit current at AC servo ampliﬁer
Capture debris in case of fracture

Table 8.9: Optional preventive measures

Acceleration sensors

Put high sensitive acceleration sensor on the unit
to detect tiny vibration
Proximity sensors
Detect extreme proximity by capacitance sensors
Contact probe sensors Locate ﬂexible contact probes in the inner hub
Visual marker
Ensure precise and easy alignment of the unit
Locating pins
Keep the unit under constraint during passing
through the inner hub
Mock-up HSC unit
Detect interference along transferring path
during the TUE work and train TUE operators

kept parallel to the inner hub. The second is applying shock absorbing pads. We adopted
NonBuren NS70, developed by Maikata-Giken, which is expected to decrease the contact
stress to 1/16 by rough estimation. The pads will be wrapped around three ﬂanges of theWFC barrel where considered to be possible contact areas. The third is employing the crash
pad at the end of the WFC barrel. The crash pad is made of thin aluminum shell insteadof Cordierite. It works to avoid the propagation of contact stress when the WFC barrel islowered in spite of inadequate alignment and mounted onto the inner hub at the end of theWFC barrel. The forth is limiting the thrust of the unit loader. The thrust is proportional
to the current for the motor, it can be controlled at the AC servo ampliﬁer of the TUE.
Finally, we need to be ready for the fracture of the WFC barrel. The ﬁfth is attaching awire cage below the inner hub during the TUE work in order to capture debris when the
fatal accident occurs.
Optional preventive measures [18, 19] are summarized in Table 8.9. Those are optional but
should be used simultaneously to reduce the risk of collision. Figure 8.27 is concept imageof preventive and post-fracture measures which are described above.

8.5.5

Proof experiment plan

We have started numerical analysis under various conditions taking suggestions from theobservatory staﬀs into consideration[20]. In addition, we are planning a proof experiment to8.5. SAFETY ISSUES ON ATTACHMENT OF HSC TO TELESCOPE

235

TUE thrust limit
limit current at
motor servo amp

Acceleration sensors
detect tiny vibration

Wedges at grip points
keep the unit parallel to inner hub

Locating pins
keep the unit under constraint

Shock absorbing pads
reduce contact stress to 1/16

Visual markers
Ensure precise alignment

Contact probe sensors
locate flexible contact sensors

Crash pad
avoid propagationof contact stress

Proximity sensors
detect extreme proximity

Wire cage
capture debris in caseof fatal accident

Figure 8.27: Preventive and post-fracture measures are labeled. Red labels are considered
to be mandatory.

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CHAPTER 8. TELESCOPE INTERFACE

Table 8.10: Experiment conﬁguration for collision case 1. Annotation is shown in Figure
8.28.
Arm of pendulum
Weight of test piece
Collision speed
Initial angle of pendulum

L
W
V
θ

1m
17.0 kg
0.06 m/s
1.1 deg

G1 flange prototype

Shock absorbing
Pad

High speed camera
Suspended G1 flange prototype (with G1 lens prototype)Figure 8.28: Plan of proof experiment. The experiment will be conducted without G1 lens
prototype.
conﬁrm whether the analysis can predict the phenomenon of fracture. In the experiment,
equivalent stress which is obtained by numerical analysis will be applied to the G1 ﬂange
prototype of the WFC barrel.
The conﬁguration for the collision case 1 is summarized in Figure 8.28 and Table 8.10. In
the collision case 1, the kinetic energy of the collision is calculated as 0.03 J. At ﬁrst, we will
start a small initial angle of pendulum, then increase the angle so that the kinetic energy
can be increased. We will conﬁrm that the prototype is still no fracture when the kinetic
energy reaches 0.03 J. The experiment for the collision case 2 will be conducted in a similarway.
We are now discussing with the manufacturers and the observatory staﬀs how to make theexperiments under control especially shape and area of the contact surface. We will decide
the experiment methods in March 2009 and carry out the experiments in April 2009.

Bibliography[1] Mitsubishi Electric Corp. 2009, HSC Report, TM-N51463[2] Morokuma, T., et al. 2008, SPIE, 7014, 161[3] Iye, M., et al. 2004, PASJ, 56, 381[4] Abazajian, K., et al. 2008, arXiv:0812.0649[5] Morokuma, T. 2009, HSC Report, HSC20090225MorokumaAG[6] Canon Inc. 1997, AG & SH
[7] Morokuma, T. 2008, HSC Report, HSC20080814MorokumaSH2[8] Mitsubishi Electric Corp. 2008, HSC Report, TM-N50940[9] Mitsubishi Electric Corp. 2008, HSC Report, TM-N50940B[10] Mitsubishi Electric Corp. 2008, Subaru HSC ﬁeld work schedule, TM-N51073-B[11] Canon Inc. 2008, HSC Report, HS40-103[12] Uraguchi, F. 2008, HSC Report, HSC20081218 UraguchiF Collision4 E[13] Uraguchi, F. 2008, HSC Report, HSC20081218 UraguchiF TUEInclination[14] Canon Inc. 2009, HSC Report, HS40-131[15] Canon Inc. 2009, HSC Report, HS40-112[16] Mitsubishi Electric Corp. 2009, HSC Report, TM-N51590[17] Canon Inc. 2009, HSC Report, HS40-120[18] Uraguchi, F. 2009, HSC Report, HSC20090128 UraguchiF SafetyMeasure[19] Mitsubishi Electric Corp. 2009, HSC Report, TM-N51640[20] Tomono, D. 2009, HSC Report, CordieriteStrength0902xx

237

Chapter 9
Data Management and Analysis
SystemThe HSC data management and analysis system is to create data products (ﬁnal images
and object catalogs) with required accuracies for HSC survey sciences. The idea describedin this document is also found in Furusawa et al. (2008)[1].

9.1

Suprime-Cam Data Analysis

9.1.1

Current situation of Suprime-Cam data analysis

Suprime-Cam is the current wide-ﬁeld camera under operation at the prime focus of Subaru
Telescope, which covers 34 × 27 square arcmin. Suprime-Cam is considered as a precursor
instrument for the HSC, since the HSC is an upgrade of Suprime-Cam in terms of its ﬁeldof view. Hence, the data set obtained with the Suprime-Cam can be used for a small testingcase of the HSC data analysis.
We have developed data analysis pipelines which compiles our experiences in processing
Suprime-Cam data frames. Figure 9.1 shows a typical analysis procedure for Suprime-Cam
data set, in which we deﬁne the data analysis for each CCD chip (bias subtraction throughsky subtraction) as ‘1a-stage analysis’, the mosaicking/stacking procedure to make a single
deep image as ‘1b-stage analysis’, and the catalog making and photometric/astrometriccalibrations as ‘2nd-stage analysis’.
The 1a-stage analysis is a series of FITS image manipulations to remove non-uniformity
in the data frames which comes from characteristics of the camera and the telescope. The
1b-stage analysis is a mosaicking procedure to make a coadded image including positionalalignment and ﬂux scaling for a set of reduced data frames. The 2nd-stage analysis is aseries of procedures to create catalogs of object parameters and to perform photometric andastrometric calibrations.
Through the development and observing operations of Suprime-Cam, we have much experience in the 1a-stage analysis and there are several dedicated software packages for thisstage We also have a robust algorithm to achieve the 1b-stage analysis (mosaicking/stacking)for Suprime-Cam data. One of the typical procedures of data analysis is implemented bynekosoft - a set of manipulation engines (Yagi et al. 2003[4]) and SDFRED - command
scripts (originally from Ouchi et al. 2003[2]).
239

240

CHAPTER 9. DATA MANAGEMENT AND ANALYSIS SYSTEM

Figure 9.1: A typical data analysis procedure for Suprime-Cam data

9.1.2

Problems in the current Suprime-Cam data analysis

The standard pipeline has successfully worked for many studies. However, there are several
problems to be studied and resolved for development of the HSC data analysis system.
First, there is no established ‘data tagging’ mechanism to measure the data quality of eachframe. The data tagging is deﬁned as a procedure to give tags on usability of the data basedon quality information which is measured by immediate analysis for frame by frame. Due tolack of the data tagging, it is diﬃcult to perform eﬃcient data selection in any later stagesof analysis. The data tagging mechanism is also important to assist eﬃcient observation
planning by feeding back the measured quality information to observers.
Second, there is no sophisticated analysis ﬂow control. The appropriate analysis ﬂow control function should assist analysis operators/users to track multiple analysis branches which
are performed for diﬀerent data set and with diﬀerent analysis parameters. Some intelligent
ﬂow control will be helpful when we perform re-analysis by adjusting input data lists andanalysis parameters to obtain the best results from a given data set. In our experience, thiskind of re-analysis is likely to happen in mosaicking/stacking and ﬂatﬁelding procedures,in which we often repeat analysis with diﬀerent input data set and diﬀerent parameters bychecking results each time we try with a particular combination of data set and parameters.
In the current Suprime-Cam analysis, most of the critical processes in the reduction procedure are manually handled on observers’ responsibility. Those processes include frame
selection for data input, checking output images, adjusting parameters for each analysis
task, and photometric and astrometric calibrations. Those manual operations in such critical parts of data analysis lead to ineﬃcient analysis, and may also lead to unreliable orun-traceable results, especially when the survey project is large and multiple people are involved in the analysis work. Since the observers maintain their own analysis logs by theirown hands, it is diﬃcult for someone else to track the analysis histories and to reproduce
and examine particular results that another researchers previously made.
Third, there are slow parts in analysis stages which could be bottlenecks in the analysis

9.2. PROTOTYPING OF HSC DATA ANALYSIS SYSTEM

241

pipelines. The slow parts comprise frequent disk I/Os for large-volume data input, especially in the 1a-stage analysis, and algorithms which have room to be reﬁned such as inmosaicking/stacking.
Those problems will be more serious for the HSC data whose size is 10 times larger thanthe Suprime-Cam data. Since the data handling by individual observers will be no longer
practical, the HSC data set should be processed to the ﬁnal products such as ﬂatﬁelded
images, mosaicked images and object catalogs in the observatory, and then provided to observers. Since the data production rate will increase up to 2GB per shot, we also have toconsider to remove any bottlenecks of the analysis pipelines and overhead times in the analysis system. Thus, in order to produce reliable processed data products in an eﬃcient way,
we try to overcome such weak points in our current pipelines and develop a new automateddata analysis system for the HSC.

9.2

Prototyping of HSC data analysis system

As a prototype of the HSC data analysis system, we start with constructing a small set ofthe data analysis system which involves essential functions to process Suprime-Cam data
currently available. The prototype analysis system is being developed on the basis of combination of existing software packages for processing imaging data and the software packages
to implement analysis ﬂow control and pipeline execution.

9.2.1

Approaches to overcome the problems in Suprime-Cam analysis

The aim of this prototyping is to address the ways of overcoming or resolving the problems
in the Suprime-Cam data analysis described in the previous section. In the prototyping, we
will introduce and test the following functions:
1. Data tagging mechanism
The data tagging mechanism will be implemented. In this mechanism, ﬁrst, real-timedata analysis system running on observing sites (Summit or Hilo) performs quality checking of each data frame by immediate data analysis after data acquisition. Here, we deﬁne
the quality information as the statistical values of seeing, noise level, sky background level,
and transparency of the sky of each data frame. Then, the data tagging determines tags
representing data usability based on the measured quality information. Thus tagged information will be used for eﬃcient frame selection in later analysis, and help us keep uniformity(mainly in terms of S/N and PSF across the survey area) of survey data in long-term survey
programs.
In order to implement the data tagging mechanism, the approach of the two-step analysis
described below (Figure 9.2) is adopted.
On-site analysis: In the ﬁrst step of the analysis, all the data frames obtained with thedata acquisition system are transferred to the HSC ﬁle server and quickly reduced with predeﬁned analysis parameters (We call this real-time analysis). The statistical values (seeing
size, noise level and sky level) on each image are measured by the real-time analysis. Thosevalues are registered in the database together with tagged information of the data usability
on a real-time basis. The information in the database is used in the second step analysis toeﬃciently select the data frames to be analyzed depending on purposes.

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CHAPTER 9. DATA MANAGEMENT AND ANALYSIS SYSTEM

Figure 9.2: The idea of the two-step analysis in the HSC analysis system: on-site and oﬀ-site
analysis.

Astrometric calibration will also be performed for each frame. If the data frame is for a
pre-deﬁned standard star ﬁeld, determination of photometric zeropoint or transparency ofthe sky will be conducted. Observers will also be able to request to make a mosaicked image
and an object catalog for a particular range of data frames taken during an observing run,in order to see the achieved depths and perform quick look for their scientiﬁc results. Thisfunction uses the tagged information in the databases.
Oﬀ-site analysis: The second step of the analysis is the post-run analysis to produce the ﬁnal
data products (mosaicked images and object catalogs) to be used for scientiﬁc analysis. The
procedure in this step includes making the best ﬂatﬁelding, performing the best calibrationsof photometric zeropoints and astrometric solutions, and creating the mosaicked images andobject catalogs with the photometric and astrometric calibrations to the precisions requiredin the scientiﬁc analysis.
Frame selection for the mosaicking analysis is performed on the basis of the tagged information on each data frame recorded by the on-site real-time analysis. Observers or data-pipeline
operators will be able to easily collect the data frames to be mosaicked onto the analysis
server which satisfy user-deﬁned criteria, such as the celestial coordinates, ﬁeld names, andother typical selection parameters (ﬁlter, exposure time, seeing size etc.), for the entire periodof an observing program, and see the achievement of the program at any stage.
2. Analysis ﬂow control system communicating with databases
For eﬃcient analysis, we aim to construct an automated data analysis pipeline whichconnects each of data analysis tasks by the ﬂow control software. The ﬂow control software
takes care of scheduling/executing/monitoring of each analysis process by deﬁning interfaces
between the analysis tasks and by handling input/output of data and memory for processingdata.
The ﬂow control software should also handle communications with databases. The databases

9.2. PROTOTYPING OF HSC DATA ANALYSIS SYSTEM

243

store information on the FITS ﬁles, information used in each analysis task including quality
and tagged information for each data frame, conﬁguration parameters for tasks, and resultson photometric and astrometric calibrations, which are necessary to perform automatedanalysis. The databases also hold all the data analysis histories (input/output data set andanalysis parameters used in each analysis), which helps to track any analysis procedures
performed by the system.
3. Improvement of data processing speed
Since the data production rate of the HSC (2GB/shot, ∼ 300GB/night) is ∼ 10 times larger
than that of Suprime-Cam, the data processing time will increase by an order of magnitude
for the 1a-stage analysis, and that for the 1b-stage analysis (mosaicking/stacking) will also
be seriously increased without improvements of data handling and algorithms of the analysis.
It is important to reduce possible bottlenecks in the analysis procedure, and the followings
are tested in the prototype system:
First, in the 1a-stage analysis, which consists of image manipulations for each CCD chip, we
divide data frames by each CCD and distribute them onto multiple analysis servers (In the
prototype, we have 6 analysis nodes with 36 CPU cores).
Second, in the analysis of the 1b-stage and the later stages, we gather a set of data frames tobe stacked into a mosaicked image and put them onto an analysis server which has multiple
CPU cores. The 1b- analysis for the mosaicked image is performed using multiple CPU
cores (In the prototype, TERAPIX software SCAMP and SWARP which have a capabilityof SMP parallel processing are used). Figure 9.3 shows the schematic view of this conceptof parallel data processing.

Figure 9.3: The idea of parallel processing of the data frames used in the prototyping data
analysis system.

In order to minimize overheads due to the disk I/O in the beginning and the end of eachanalysis task, we also test on-memory analysis without frequent disk I/Os in the prototype

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CHAPTER 9. DATA MANAGEMENT AND ANALYSIS SYSTEM

system (Figure 9.4).

Analysis Flow Control Softwaremodules
on a pipeline

readfits

overscanastrometircsubcalibration

FITS

WCS

files

solution

writefits

FITS
files

Figure 9.4: An example of analysis by the analysis ﬂow control software, which makes ananalysis pipeline by connecting analysis tasks (data processing engines) and handles the
input/output of FITS data. This ﬁgure also shows the idea of on-memory data processing
by the analysis ﬂow control software.

9.2.2

Development steps

The ﬁrst prototype analysis system will be examined with a small set of testing data andtested in the actual observing runs. The system will also be applied to a wide variety ofmassive data set in the data archive.
Through this application, ﬁrst, we will evaluate the capability of the analysis framework
software, then decide whether we can proceed with the software for further implementation of the HSC data analysis system, or we need to ﬁnd or develop another framework byourselves. Second, we will evaluate our implementation of the analysis pipeline including algorithms, information registered in the databases, and ﬁnd potential problems in processing
the upcoming huge HSC data set. We will investigate the system optimization by measuring system performance and loads on individual hardware resources, and discussing which
process should be distributed or parallelized.
In addition to testing with the Suprime-Cam data set, we will also develop artiﬁcial data
simulating HSC data in order to examine our analysis pipelines and to know what parts ofthe pipelines are available to the HSC data set and what parts should be modiﬁed. The
current plan of developing steps is listed below.
Phase 1: testing of the on-site analysis for Suprime-Cam data (present-2009.9)
Phase 2: development and testing of the oﬀ-site analysis for Suprime-Cam data (2009.102010.3)
Phase 3: development and testing of the on-site analysis for HSC with simulated data foreach CCD (2010.4-2010.12)
Phase 4: development and testing of the oﬀ-site analysis for HSC with sets of simulated
mosaic data (in 2011)
Once the prototype system is evaluated, we will scale up the system to the full HSC-ANA
system based on the evaluation results.

9.2. PROTOTYPING OF HSC DATA ANALYSIS SYSTEM

9.2.3

245

Goals of prototyping phase 1 and 2

In the prototyping phase 1 and 2, we aim to construct the analysis system which provides“nightly reduced data” to observers, and assists them to make their exposure plan using
the results derived by on-site analysis. Also, the prototype system will enable observers to
perform quick analysis of their data set during the night. By oﬀering the prototype system
to open-use observers, we will obtain feedbacks from them which will be important to discuss
the ﬁnal design of the HSC data analysis system.

9.2.4

Current status of the phase 1 prototype

We are now developing the phase-1 prototyping system, in which we adopt the following
software components.
1. Data processing engines (Analysis applications)
Most of the applications for the 1a-stage analysis are developed by ourselves based on C
and C++ languages using the CFITSIO library (Pence 1999[3]) for manipulations of FITS
ﬁles. For the statistical value measurement and astrometric/photometric analysis and the
1b-stage analysis, software derived from the TERAPIX project (http://terapix.iap.fr/)(SExtractor v2.5.0, SCAMP version 1.4.4-MP, SWarp version 2.17.1-MP). The TERAPIX
software is a good starting point for us to quickly construct the prototype system, especiallyfor the 1b- and 2nd-stage analysis, since it has been well tested by astronomical communities
including CFHT Legacy Survey and frequently improved. It is also under GPL license andwe can easily modify the source codes when needed. Importantly, the above TERAPIX
software has a SMP mode, which allows to easily run the threaded programs on multiple
CPU cores, which is particularly useful in the 1b-stage analysis.
2. Analysis ﬂow control software
In the prototype, we use a combination of two analysis ﬂow control software packages toconstruct eﬃcient analysis pipelines.
We use a software called “BASF (Belle AnalysiS Framework)” developed by KEK (High
Energy Accelerator Research Organization in Japan) which was originally developed forprocessing data with the very high data rate in particle physics experiments at the B-Factory
facility at KEK. The BASF connects multiple analysis modules (data processing engines) to
ﬂexibly construct a series of automated analysis pipelines, and handles data input/output
between the analysis modules. In this prototype, we adopt an upgraded version of BASF,“Roobasf” which is combined with another analysis framework ROOT developed by CERN.
ROOT adds to the original BASF a strong capability of handling any input data to beprocessed as objects on a memory space without disk I/O. The Roobasf ﬁrst reads FITS
ﬁles, and loads the FITS header and pixel data onto the memory space. Then it processes
the data on the memory without writing anything out to the disk, and ﬁnally writes resulting
FITS ﬁles onto the disk after a series of the analysis pipelines are completed.
Scheduling, execution, and monitoring of the data analysis pipeline are handled by another ﬂow control software - RCM (R&D Chain Management) developed by a commercial
vendor Quatre-i Science Corporation1 . The RCM has a capability of constructing analysis
pipelines and providing them with interfaces communicating with XML-based databases.
The databases can ﬂexibly record data structure with the aid of XML structuring. This1

http://www.i4s.co.jp

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CHAPTER 9. DATA MANAGEMENT AND ANALYSIS SYSTEM

capability allows us to easily maintain complex analysis histories and easily add newly deﬁned information without diﬃculty modifying schema, which leads to usually a problem in a
typical relational database. The RCM is an XML-based system and provides user interfaces
which can be accessed through any XML browsers. The RCM also has a function to conduct
process distribution to multiple analysis servers. This function is applied for the 1a-stage
analysis in our phase-1 prototype system.
3. Databases
In order to realize ﬂexible recording of various information and to identify necessary metadata for our system, we use two diﬀerent kinds of databases – relational database (RDB;
PostgreSQL version 8.3) and XML-based database which is a component of the RCM software. Basically, simple meta data including the FITS keywords are registered in the RDB,
and the other information generated by the analysis pipeline including statistical values
measured by on-site analysis and any analysis histories are registered into the XML-based
database. The RCM is responsible for communication between the analysis processes andboth databases.

Fits raw data

File lists

Logparameters
Relational DB
Analysis Flow Control, DataBase I/F, User I/F

Fits header
File location(RCM)
processXML DB

distribution

Analysis history
Analysis Flow Control

Analysis parameters(Roobasf)

QA/QL parameters

Realtime

Pipeline 1a

E OSR subtraction
 FlatfieldingE AG probe masking
E
1chip analysis
results

On demand

Pipeline 1b

E Mosaicking
 Stacking
E
Pipeline 2

E Catalog making

E Statistical values(FWHM, Sky level, # of Stars)
 AstrometryE Standard star ana
EData Processing Engines(Applications)Figure 9.5: System diagram of the current prototype system for Suprime-Cam on-site analysis (phase 1), which is composed of RCM ﬂow control software and Roobasf analysis framework.

A schematic view of the current system is shown in Figure 9.5. The hardware system used
for the prototyping phase 1 and 2 (see Figure 9.6), on which the above software components
are operating, is as follows:
1. ANA and HANA – analysis server as a component of observing control system (OCS).
All the data are transferred from the Suprime-Cam data acquisition PC to this node as soon
as the data is obtained. The raw data are then transferred to the below analysis servers in
RCM.
2. Analysis servers – type 1× 2 nodes (CPU:Intel Xeon QuadCore 2.6GHz ×1, Memory:2GB,
HDD:500GB), type 2× 3 nodes (CPU:Intel Xeon QuadCore 2.6GHz ×2, Memory:4GB,
HDD:720GB). All the data analysis processes are executed. Roobasf runs as a process

9.2. PROTOTYPING OF HSC DATA ANALYSIS SYSTEM

247

on these nodes when invoked by the RCM control server. Data is arrived on each of theservers on this nodes.
3. Database server – CPU:AMD Opteron 2.8GHz ×2, Memory:16GB, HDD:250GB. This
server is in charge of managing all the databases for the RCM analysis system.
4. Web/Control server – CPU:Intel Xeon QuadCore 2.6GHz ×1, Memory:2GB, HDD:500GB.
This server is in charge of analysis ﬂow control, communicating with databases, and distributing processes onto analysis servers. This server also provides the user interface for operating
the RCM analysis system.
Snapshots of the current implementation of the prototype system phase 1 based on the
RCM+Roobasf framework are shown in Figure 9.7. Currently, the real-time part of theon-site analysis (1a- and 2nd-stage analysis) has been implemented, except for the noise
measurement of the data frame and the standard star analysis (See Figure 9.5). The realtime system can obtain seeing size and sky background level for each frame and registers thosestatistical values to the databases. We are working to deﬁne the useful tagged information
based on those values. The on-demand 1b-stage analysis has also been implemented, inwhich analysis users can select data frames using information in the databases recordedin the real-time analysis. The catalog making part is now under development. We are also
working on modiﬁcation of the user interface for eﬃcient use of the prototype system, aiming
to provide the system to visiting observers.

Figure 9.6: The hardware system used for the prototyping phase 1 and 2.
After the ﬁrst trial of phase-1 prototype system, we will next develop the oﬀ-site analysis
system (phase-2 prototype) for Suprime-Cam data and test the following things:(1) We will test utilization of the tagged information to process mosaic data set from multiple
observing runs. We aim to establish a way of photometric and astrometric calibrations for248

CHAPTER 9. DATA MANAGEMENT AND ANALYSIS SYSTEM

given observing data. We will also evaluate the prototype system by providing the system
to open-use observers.
(2) For better analysis ﬂow control, we will study a way of examining analysis results and try
to implement automation of re-analysis by adjusting analysis parameters to obtain reasonable
analysis results. The system is also examined for the management of analysis branches fordiﬀerent re-analysis trials.
(3) For faster observing speed, we will test the system performance and study a way ofscaling up the prototyping system to the HSC system.

9.3

Key Issues of HSC Data Management and Analysis
System

9.3.1

Diﬀerences between Suprime-Cam data and HSC data

The signiﬁcant diﬀerences between the Suprime-Cam and HSC data, which have to be addressed to scale up the prototyping system to the HSC analysis system, are shown below.
(1) Larger data size
The data size is 10 times larger and the covered area is 7 times wider than those of SuprimeCam data. In addition, the survey coverage in the HSC era is supposed to be up to ∼ 1000
deg2 , which is hundred times wider than those obtained by Suprime-Cam. Thus, the HSC
data set requires better analysis eﬃciency and faster analysis speed.
(2) Larger Field-Of-View
The large ﬁeld of view of HSC (φ1.5 degree) will bring us diﬃculty in mapping the survey
ﬁeld with uniform signal-to-noise ratios (S/Ns) and in making mosaicking/stacking images,
and even making object catalogs in a simple manner as in the Suprime-Cam analysis. Thosechange requires adequate pointing strategy to make uniform S/Ns across the survey ﬁeld
as much as possible. Also, an appropriate algorithms to make mosaicking/stacking andobject extraction have to be developed. We also need careful handling of ﬂatﬁelding by
understanding the characteristics of the ﬂat illumination pattern of the HSC.
(3) Larger geometrical eﬀects
We expect strong distortion of the corrector optics and geometrical eﬀects in images due
to atmospheric dispersion corrector (ADC). Those geometrical eﬀects will change against
telescope pointing, especially in elevation. This requires adequate treatments of geometricaltransformations and resampling of pixel data in our analysis processes, especially for better
ﬂatﬁelding, mosaicking/stacking, and catalog making.

9.3.2

Challenges to scale up Suprime-Cam prototyping system

In order to scale up the prototype system, we have to discuss and address the following
points in addition to improvements of the problems described in Section 9.1.2.
(1) To achieve required accuracies of data products:
1. Adequate algorithms in image reduction has to be developed and well tested. We have toﬁnd reliable methods for ﬂatﬁelding, distortion/ADC eﬀect correction, and sky subtraction,by using simulation data and engineering observations. In particular, it is important toadjust and establish those algorithms using data from engineering observations.

9.4. GOALS OF HSC DATA ANALYSIS SYSTEM

249

2. The ways of photometric and astrometric calibrations need to be discussed. We may
determine optimized observing strategy of how we take reference data. Necessary external
data for calibrations have to be determined (standard star network etc.)
3. We have to investigate optimal procedures and algorithms of mosaicking, stacking, object
detection, and extraction of object parameters (photometry, shape). This should be discussed in combination with an optimized strategy of pointing and exposures in the typical
HSC survey observations.
(2) To improve analysis eﬃciency and analysis speed:
1. We need to optimize information to be stored in the databases and analysis ﬂow. Thisshould be investigated during the prototyping.
2. The system design of hardware (CPUs, disks, network) for the HSC data analysis system
will be determined based on the system performance measured by the prototype system.
We will study the optimal hardware conﬁguration in conjunction with analysis algorithms,
especially for mosaicking/stacking procedure.
(3) To evaluate analysis system:
1. We need to develop artiﬁcial data simulating HSC data to examine the HSC data analysis
system. Developed software has to be tested in each algorithm and in each component. Basic
test for the analysis software will be done with the simulated data to ﬁnd distinct problems.
Finalizing of the software will be done with the engineering data.
(4) To maximize science output:
1. We need to develop a science archive system for mosaic image and catalog data of the
HSC. We will study what information has to be available in the science archive and what isa reasonable interface for the eﬃcient science use of the HSC data products.

9.4

Goals of HSC Data Analysis System

9.4.1

Data products

The targeted data products in the HSC data analysis system are summarized below. The
valus in the parenthesis are target accuracies in the HSC data analysis tentatively set, someof which should be disucussed with science work groups.
(1) Nightly Products (by On-site Analysis)
1. Observing logs with quality assurance information
2. Preliminary ﬂat (systematic error: peak-to-peak < 5%)
3. Reduced images for individual CCDs with calibration
4. Best-eﬀort Mosaicked images with calibration (alignment error: < 1 pix rms, ﬂux error:<
10% rms)
5. Shallow object catalogs (errors in ﬂux: < 10%, astromtry: < 1arcsec rms, object parameters: TBD)
(2) Post-run Products (by Oﬀ-site Analysis)
1. Best ﬂats (systematic error: peak-to-peak < 2%)
2. Deep mosaicked images with better calibration (alignment error: < 0.7 pix rms, ﬂux
error: < 5% rms)
3. Noise or exposure weight map
4. Deep object catalogs (errors in ﬂux: < 5%, astrometry: < 0.5arcsec rms, object parameters: TBD)

250

9.4.2

CHAPTER 9. DATA MANAGEMENT AND ANALYSIS SYSTEM

Analysis time

The targeted analysis time is currently set as follows:(1) To avoid waiting time in on-site analysis for the 2nd exposures, the analysis for 1 shot
has to be done in ∼ 25sec after the data arrive in the HSC analysis system.
(2) To obtain analysis results in a reasonable time, the data transfer time from the DAQ
machine to the HSC analysis system should be much shorter than 35sec.
(3) To catch up with the nightly data rate by processing the data set into mosaic images and catalogs, processing rate for mosaicking/catalog making has to be larger than
700MB/minute.

9.5

Summary

The data analysis procedure established for a precursor instrument Suprime-Cam has beensuccessfully used. However it has problems to be addressed for HSC analysis. We havestarted with prototyping based on the Suprime-Cam analysis procedure to test methods toovercome the problems. Studies for scaling up the prototyping system has been started. Requirements of HSC data management and analysis system are under discussion with sciencegroups.

9.5. SUMMARY

251Òshow detailÓ button opens ÒMonitor windowÓ

Monitor window

Monitor window

If you want check fits in detail,Links under the fits thumbnail opens d9.
(you should start X server on your local computer
before you push this link)

Figure 9.7: Snapshots of Suprime-Cam on-site analysis (phase 1) which is composed of RCM
ﬂow control software and Roobasf analysis framework. The top panel shows a window foranalysis control for the on-site analysis and the summary window of a processed SuprimeCam data. The bottom shows the window for the detailed information of the processedshot.

Bibliography[1] Furusawa, H. et al. 2008, SPIE, 7016, 45[2] Ouchi, M. et al. 2003, PhD thesis, University of Tokyo[3] Pence, W. et al. 1999, ADASS VIII, 172, 487[4] Yagi, M. et al. 2003, AJ, 123, 66

253

Chapter 10
HSC Project Schedule
Schedule of the Hyper Suprime-Cam project is summarized in Table 10.1 and Table 10.2.
All of the camera components are being designed and prototyping of some of components
are currently underway such as shutters and ﬁlters. Manufacturing and testing of these aredue in March 31, 2010. All the CCD procurements and testing should be completed by thesame dates. Then, we start the assembly of the camera from April 2010.
The PFU and the WFC is due in August 31, 2010 and delivered to Mitaka. Assembly of
the HSC entire system start in the following week and the system test will be carried out
until the end of the year. We ship the system to Hawaii in January 2011 and we reserve 3
months for the shipment by keeping some spare time there. Although we plan to assemble
the system at Advanced Technology Center at Mitaka right now, all the necessary condition
might not be met such as crane availability. A factory of Mitsubishi is our backup plan forthe place of the assembly.
Part of the telescope modiﬁcation works will start in April 2010 and the top unit modiﬁcation work will be done by overlapping with the shutdown time of the 8.2 m M1 coatingin August 2010. Re-routing of the cables and installation of the clipping mechanism isscheduled in June 2011.
After the arrival of the HSC system in Hawaii, we carry out the ﬁnal check of the systemat the telescope till the end of August 2010. The function test of the PFU during the day
time is scheduled till September 12, 2011. Right now PFU is planned to be mounted on thetelescope during the testing. This means that telescope cannot be used during night it ishard to switch the instruments so quickly. This is completely the waste of telescope time andwe should work out how to test PFU located at the PFU ﬂoor independent of the telescope.
We hope we start the commissioning run from September and it will take for roughly amonth to test all the functionarity. During the commissioning we expect to unloading andloading of HSC system several times because minor modiﬁcation and tuning, such as shim
insertions, will be required based on the test results at the telescope. It will be better if thecommissioning run is split into several phases.

255

10.3.31

10.3.3110.3.3110.3.31

09.4.109.4.109.4.1

Dewar (design and manufacturing)

Shutter

Filter EX

Electronics

CCD (procurement and10.3.3110.7.30

10.8.3110.8.3110.8.3111.7.27

10.4.1

10.8.209.4.109.4.1

10.5.14

AssemblyTest

10.6.30

10.9.29

11.7.27

10.7.1

11.7.1Top Unit Modification

Cabling & clip mech. installation

11.10.26

10.12.31

11.3.3110.9.111.1.3

Assembly and test

Shipping to Hawaii

Camera test

11.10.26

11.7.20

11.9.13

PFU test (daytime work)

Commissioing Observing Run

2009

2009

2010

2011

2011

2010/09 HSC System Assembly Start

2010/04 Camera Assembly Start

2010

Figure 10.1: HSC Project Schedule. Telescope cannot be used during the red designated period.

11.8.1

11.9.1211.4.1HSC System

10.9.1

Preparation

10.5.14

Telescope Modification

Design, Manufacturing & AssemblyPrime Focus Unit

Wide Field Corrector

Manufacturing

Camera Assembly and test

10.8.3110.8.3110.4.1

Filter (procurements)

09.4.1

11.9.30

09.4.109.4.1

10.3.31

09.4.1test)

11.9.30

Camera09.4.1Finish Date

Start Date

09.4.1

Activity Name

256
CHAPTER 10. HSC PROJECT SCHEDULE

11.10.31

09.9.30

10.3.31

09.4.109.10.1P2 prototype (off-site for SC)

P3 prototype (HSC simulation for11.10.31

11.10.31

11.9.1

10.4.1

Testing with engineering data

Science archive

mosaic)P4 prototype (HsC simulation for

single CCD)
11.8.3110.3.3110.12.31

10.4.111.1.310.3.3109.10.1

System desing for HSC prototype

Study of application Algorithms

09.4.1

Data Analysis and Management

09.4.1Finish Date

Start Date

P1 prototype (on -ste for SC)

Activity Name

2010

2010

Figure 10.2: HSC Project Schedule (continued)

2009

2009

2011

2011

257

Appendix A
Notes for observers
While we were desgining HSC, we noticed several issues that are diﬀerenct from SuprimeCam. In this appendix, we list up those information for future the HSC observer’s convenience.
1. Because of the atmospheric dispersion and the asymmetrically variable distotion parameter (see Chapter 4), image is smeared to some extent during the exposure.
HSC20080930 KawanomotoS MaximumExposureTime discussed the possible
length of the exposure time. The degree of smear also depends on which direction you
look at.
2. HSC20081028 KawanomotoS ADCMaxExp discussed impact of the image shift
and the asymmetric distortion caused by the ADC in the WFC. This document concluded that this additional distortion hardly has bad inﬂuence upon the image quality
under the condition that the auto-guider is operating.

259

ケーブル巻取り装置
インスツルメント ロテータ モータ

ジャッキトップ フレーム
シンクロユニット

クロスローラ ベアリング

ヘキサポッドフィルタ交換機構

デュワー
インターフェイス
フレーム
CCDデュワー

CCD
ロテータ エンコーダ
158 cmフィルター

シャッター
レンズ インターフェイス
フレーム
ベース フレーム
カービック カップリング

⼤気分散補正レンズ

補正光学系