WIDE FIELD FIBER-FED OPTICAL MULTI-OBJECT SPECTROMETERWFMOSSTUDY SUMMARYJANUARY 30, 2009

Wide Field Fiber-Fed Optical Multi-Object
Spectrometer (WFMOS)SCIENCEWFMOS is the premier tool for addressing thenature of dark energy and the assembly history ofgalaxies such as M31 and our Milky Way.
WFMOS will:• Provide a constraint on the dark energy equation ofstate, w, to 3%

Dark Energy

Low Redshift

High Redshift

Redshift Range
Number of galaxies
Spectral Range

0.7 ≤ z ≤ 1.6> 4,000,000
6300–9700 Å

2.5 ≤ z ≤ 3.3> 100,000
4200 – 6500 ÅSpectral Resolution

3500

1500

RAB < 23.0

RAB < 25

Low ResHigh Res

Magnitude Limit• Measure the growth rate of density fluctuations, fg , to1.5% at 0.7<z<1.6 and 7% at 2.5<z<3.3• Measure the mass, extent, and substructure of the Milky
Way dark matter halo

Galactic
Archaeology• Conduct kinematic measurements of dwarf spheroid
candidates

Number of Stars

3 x 106

1 x 106

Spectral Range

8150–8850 Å
4800–5500 Å

4800–6800 Å• Measure the extent, substructure, and kinematics ofMilky Way stellar populations

Spectral Resolution

5000

20,000

V=20,
S:N10-15/Å

V=17
S:N 150 / Å• Explore the fossil record of chemical evolution of stellar
populations in the Milky Way

Sensitivity Limit

INSTRUMENT
The instrument survey efficiency is maximized with highly
accurate positioners, fast reconfiguration time, broad
spectral range, and high quantum efficiency.

Rigid Fiber Positioners that reconfigure in ~40 sec
with 5 µm translational accuracy and no tilt
Prime Focus Instrument

One of three identical Schmidt
Spectrographs offering flexible spectral
resolution and wavelength coverage

Detectors with high quantum efficiency in the blue andred range

Each of the 3 Spectrographs

Detectors

Collimator f-ratio

f/2.12Detectors / spectrograph2

Camera f-ratio

f/1.07Detector format

Minimum beam size~300 mm diameter

93 mm x 46.5 mm
5760 x 2880 pixels

Fiber pitch at slit

230 µm, center-center

Pixel size

16 µm

Projected fiber size at detector

54 µm diameterRed optimization

450 µm thick fully
depleted p-channel

SRE* width at detector

54.0 µm

Blue optimization

Delta-doped

Projected fiber pitch

116.25 µm

AR coating

custom dual layer
TiO2/SiO2* SRE – Spectral Resolution Element

MANAGEMENT
Development of this ground breaking instrument with an international team is accomplished with amanagement strategy developed from a history of delivering complex science instruments.
Partner

Responsibility

Cost $K

JPL

Mgmnt, Sys Eng.,
Positioner, PFI, A&G,
I&T

27,237

Cambridge

Spectrograph

10,383

PSU

Positioner Software

563

Positioner/ PFI

CaltechDetector System

2,578

Fiber System

LNA

Fiber System

3,550

UK-ATC
UCL

System Software, Data
Reduction
Metrology System

2009

2010

2011

PHASE B
Kickoff 1/15

8,000

Baseline Total

58,305Reserves

10,222

Total w/ Reserves

68,527

8/1

2014

PHASE D

Project
CDR

9/12
Integration & TRR

Positioner
EM SysE

12/12
PSR/SIR

Del Bench

Fiber Cable A
Cable B

4,355

Subaru Elements

2013

PHASE C
Project
7/1 PDR

Detector/
Spectrograph System

1,639

2012

System Software
Instrument
Integration@ JPL

MGSE
Avail

Detector
Sys 1

Del Prelim
Sys Software

Del PFI Assembly
to I&T

Rec PFI

Det 2

Det 3

PSR

Ver2
Recipes

Cable A

Ver3
RecipesSpectrograph
Del Sys SW
Ship

CamerasSpectographArriveIntegration @ Subaru

Del
WFMOS
7/11

WIDE FIELD FIBER-FED OPTICAL
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SUMMARY REPORT

Contributors to the WFMOS Study:
Richard Ellis, Principal Investigator, Caltech
Mary White, Project Manager, JPL
Robin Bruno, Proposal Manager, JPL
Michael Seiffert, Project Scientist, JPL
Stuart Lynn, IfA, Univ. of Edinburgh
Peder Norberg, IfA, Univ. of Edinburgh
John Peacock, IfA, Univ. of Edinburgh
Fergus Simpson, IfA, Univ. of Edinburgh
Masahiro Takada, IPMU, University ofTokyo
Masami Ouchi, Carnegie Observatories
Thomas Kitching, Univ. of Oxford
Filipe Abdalla, Univ. College London
Ofer Lehav, Univ. College London
Ignacio Ferreras, MSSL—Univ College
LondonLaerte Sodre, Sau Paulo
Scott Chapman, Inst. of Astronomy, Univ.
of Cambridge
Mike Irwin, Inst. of Astronomy, Univ. ofCambridge
Geraint Lewis, School of Physics, Univ. of
SydneyRodrigo Ibata, Observatoire de Strasbourg
Vanessa Hill, Observatoire de Paris
Alan McConnachie, Dept. of Physics andAstronomy, Univ. of Victoria
Kim Venn, Dept. of Physics and Astronomy, Univ. of Victoria
Mark Wilkinson, Dept. of Physics & Astronomy, University of Leicester
Nobuo Arimoto, National Astronomical
Observatory of Japan
Amina Helmi, Kapteyn AstronomicalInstitute, Univ. of Groningen

Chris Evans, UK Astronomy Technology Centre, Royal Observatory, Edinburgh
Beatriz Barbuy, Universidade de Sao Paulo,
IAG
Luciana Pompeia, IP&D, Universidade doVale do Paraiba
Dante Minniti, Dept. of Astronomy, Pontificia
Universidad Catolica
Rich Dekany, Caltech
Anna Moore, Caltech
Roger Smith, Caltech
Steven Beard, UK Astronomy Technology
CentreIan Bryson, UK Astronomy Technology Centre
Andy Vick, UK Astronomy Technology CentrePeter Doel, Univ. College London
Dave King, Univ. of Cambridge
Ian Parry, Univ. of Cambridge
Dave Braun, JPLCharles Fisher, JPL
Amanda Frieze, JPL
Larry Hovland, JPL
Muthu Jeganathan, JPL
Joel Kaluzny, JPL
Roger Lee, JPL
Norman Page, JPL
Lewis Roberts, JPL
S. Tere Smith, JPL
Ron Steinkraus, JPL
Karl Reichard, Penn State University
Antonio Cesar de Oliveira, LNA, Brazil
Ligia Souza de Oliveira, LNA, Brazili

WIDE FIELD FIBER-FED OPTICAL
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Table of Contents
1.0 INTRODUCTION ...................................................................................................... 3
1.1
Proposal Organization .............................................................................................5
2.0 SCIENTIFIC CASE AND REQUIREMENTS............................................................. 6
2.1
Context.......................................................................................................................6
2.2
Dark Energy and Modified Gravity.........................................................................7
2.2.1 The Challenge of Cosmic Acceleration ....................................................7
2.2.2 A Dark Energy Strategy ..............................................................................7
2.2.3 Baryonic Acoustic Oscillations ...................................................................7
2.2.4 Modified Gravity and Redshift Space Distortions ...................................8
2.3
Proposed WFMOS Surveys ...................................................................................92.3.1 Survey Strategy: Changes Since the Feasibility Study..........................92.3.2 Synergy with HyperSuprimeCam ............................................................11
2.4
Galactic Archaeology.............................................................................................122.4.1 Near-Field Cosmology Comes of Age ....................................................122.4.2 Survey Strategy: Differences with the Feasibility Study ......................13
2.5
Additional Science..................................................................................................142.5.1 Determining the Neutrino Mass ...............................................................142.5.2 Studies in Galaxy Formation ....................................................................15
2.6
Conclusion...............................................................................................................15
3.0 INSTRUMENT DESCRIPTION............................................................................... 18
3.1
Rationale..................................................................................................................18
3.2
Instrument Overview..............................................................................................18
3.3
Major Requirements ..............................................................................................22
3.4
High-Level Assumptions .......................................................................................25
3.5
Instrument Trades ..................................................................................................27
3.6
Metrology Camera System ...................................................................................27
3.7
Field Element ..........................................................................................................29
3.8
Acquisition and Guide Camera ............................................................................30
3.9
Prime Focus Instrument ........................................................................................31
3.10 Positioner.................................................................................................................34
3.11 Fiber System ...........................................................................................................42
3.12 Spectrograph...........................................................................................................45
3.13 Detector System .....................................................................................................50
3.14 System Software ....................................................................................................53
3.14.1 Observation Preparation ...........................................................................53
3.14.2 Observation Management ........................................................................53
3.14.3 Observation Execution ..............................................................................54
3.14.4 Engineering and Test Software ...............................................................54
3.15 Data Reduction .......................................................................................................56

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4.0 INSTRUMENT PERFORMANCE ........................................................................... 58
4.1
Optical Throughput ................................................................................................58
4.2
Reconfiguration Time ............................................................................................60
4.3
Allocation Efficiency ...............................................................................................60
4.4
Exposure Times and Survey Time Estimate......................................................61
4.5
System Versatility and Upgrades ........................................................................61
4.5.1 Positioner.....................................................................................................61
4.5.2 Spectrograph ..............................................................................................61
5.0 INTEGRATION AND TEST .................................................................................... 63
5.1
Integration Approach .............................................................................................63
5.1.1 Spectrograph/Detector/Data Pipeline/WFMOS Control Subsystems’Path ..............................................................................................................65
5.1.2 Prime Focus Instrument (PFI) Path.........................................................65
5.1.3 Integration and Verification at Subaru Telescope.................................66
6.0 OPERATIONS ........................................................................................................ 68
6.1
Data Reduction Package B ..................................................................................68
6.2
Conducting the Dark Energy and Galactic Archaeology Science Surveys...68
6.3
Installation and Removal.......................................................................................69
6.4
Additional Considerations .....................................................................................69
7.0 MANAGEMENT AND COST .................................................................................. 70
7.1Management and System Engineering ..............................................................70
7.2
Instrument Cost ......................................................................................................72
7.3
Schedule..................................................................................................................76
7.4
Risk Management ..................................................................................................79
7.5
Reserves..................................................................................................................81
8.0 REFERENCES ....................................................................................................... 82
9.0 ACRONYMS AND ABBREVIATIONS .................................................................... 83
Table of Figures
Figure 1.0-1: WFMOS Instrument System Elements ............................................................ 4
Figure 1.1-1: Organization of submittal. Our concept study is summarized in the Study
Summary (this document). Separate volumes include: Management and Cost Volume,
Detailed Technical Design, the Operational Concept Design Document, and the
Functional Performance Requirements Document................................................................. 5
Figure 2.2-1: (Left) The redshift-space correlation function ξ (rp, π) at <z> = 0.8, shownin false color to emphasize the Kaiser anisotropy at amplitudes ξ = 0.1 from the recent
survey of Guzzo et al. (2008). The ‘finger of God’ radial stretching arises from virialized
random motions at separations of a few Mpc, indicating the degree of dark matterclustering at that epoch. (Right) The growth rate function fg(z) inferred from redshiftspace distortions using various existing and planned surveys (black/blue points) and theprecision that can be expected from a WFMOS survey of 4 ×106 galaxies over 0.7 < z <1.6 and 100,000 galaxies at z ~3 (red points, each split into two redshift bins). See
OCDD for discussion of the theoretical curves. ...................................................................... 9iii

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Figure 2.3-1: Fractional precision of various surveys. Left panel illustrates the expected
fractional precision of various surveys in measuring the acoustic scale (using our own
S/N formula, as discussed in the OCDD). Right panel shows the fractional precision in
the growth rate (following White, Song & Percival 2008). The observed results for the
2dFGRS and SDSS at low redshift are plotted as points, together with the predictedaccuracy of the VIPERS VLT survey. The lines are intended to show the precision
expected in redshift shells of width 0.1. In the left-hand panel, we also show (triangle
points connected by dotted lines) the average precision expected from the full volume of
BOSS and of our WFMOS surveys......................................................................................... 11
Figure 2.4-1: A measure of the impressive data-gathering power of WFMOS incomparison to other wide field spectroscopic facilities. The ordinate represents the
product of the multiplex and number of resolved spectral elements above a fixedsignal/noise per unit observing time. Our team has learned much from its extensive useof the FLAMES and DEIMOS instruments, and this has been reflected in our WFMOS
requirements. .............................................................................................................................. 13
Figure 3.2-1: WFMOS Block Diagram ................................................................................... 19
Figure 3.6-1: Schematic of Single Camera on the Prime Focus Support Strut .............. 28
Figure 3.6-2: 1100 Series Camera from Spectral Instruments, Inc. ................................. 29
Figure 3.6-3: Optical Layout of the Subaru Telescope and the Metrology Camera....... 29
Figure 3.7-1: The WFMOS Field Element–to–Fiber Tip Coupling Efficiency .................. 30
Figure 3.8-1: The proposed configuration of the WFMOS A&G System. A coherent fiber
bundle reimages the focal plane onto the A&G Camera. .................................................... 31
Figure 3.9-1: The Prime Focus Instrument ........................................................................... 32
Figure 3.9-2: The Rotating Portion of the PFI ...................................................................... 33
Figure 3.9-3: The PFI Fixed Structure ................................................................................... 33
Figure 3.10-1. Positioner Bench Populated with One Module and a Single Positioner . 35
Figure 3.10-2: Positioner Patrol Region Overlap. The black dashed lines indicate thepatrol regions of adjacent positioners, showing there is complete coverage of the
populated field of view. The distance between adjacent positioners, d, is 8 mm, while
the diameter of the Patrol Region, 2R, is 9.5 mm. Blue and orange represent the twostages of the positioner motors (see text).............................................................................. 36
Figure 3.10-3: Optical coupling efficiency versus distance from the center of the field ofview. The plot shows the field coverage of the positioner array. Coupling includes the
optical performance of the WFC and the numerical aperture of the fiber, but does notinclude seeing or WFC internal flexure................................................................................... 37
Figure 3.10-4: The Prototype Cobra Fiber Positioner ......................................................... 38
Figure 3.10-5: Rotary Squiggle Motor Motive Force Schematic........................................ 38
Figure 3.10-6: Patrol Region of a single fiber positioner (top view). The two motors ineach positioner have axes that are offset from one another, allowing the fiber tip totravel anywhere within the patrol region................................................................................. 39
Figure 3.10-7: Positioner System Block Diagram ................................................................ 40
Figure 3.10-8: Prototype Positioner Testbed. A light source illuminates a fiber optic
mounted in the Cobra positioner unit under test. A CCD camera accurately measures
the position of the fiber tip......................................................................................................... 42
Figure 3.11-1: Schematic of the fiber system from the positioners to the spectrograph.
Pictures of key hardware from risk-reduction activities are also shown............................ 43
Figure 3.11-2: Fiber Throughput and the Contribution from Various Components ........ 44
Figure 3.12-1: The optical design of a single WFMOS spectrograph, shown here asconfigured for the GAHR survey. The diameters of the major optical elements areshown in millimeters. ................................................................................................................. 48
Figure 3.12-2: CAD model of one of the three identical spectrographs, shown here asconfigured for the GAHR survey. The table is approximately 5 m × 3 m.......................... 48

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Figure 3.12-3: A time sequence showing the system reconfiguring between the DELZ
mode (top left) and the GAHR mode (bottom right). Note that camera, the filter wheel
and the grating exchanger all move at the same time. ........................................................ 50
Figure 3.13-1: The CCDs will be located within an evacuated enclosure supported by anarrow duct carrying a cold finger from a Joule Thomson (low vibration) closed-cycle
cooler outside the spectrograph beam. .................................................................................. 51
Figure 3.13-2: Theoretical 1-R and lower bound on QE compared to next bestcompetitor. The Hamamatsu can be improved in the red by increasing thickness but
only to 300μm, and by applying the Ti02/SiO2 coating designed by JPL. ......................... 52
Figure 3.13-3: Read noise versus dwell time for three DECam CCDs manufactured by
LBNL/DALSA. Courtesy of Juan Estrada, Fermi Lab........................................................... 53
Figure 3.14-1: Control Hierarchy ............................................................................................ 55
Figure 3.14-2: Block Diagram. Subaru infrastructure is shown in green, Subarusupplied components are shown in cyan, and WFMOS-supplied components are shownin yellow. ...................................................................................................................................... 55
Figure 3.15-1: DR Subsystem Block Diagram and Interfacing .......................................... 56
Figure 3.15-2: DR Data Flow .................................................................................................. 57
Figure 5.1-1: The Instrument Integration and Test flow follows two parallel paths before
final integration onto the telescope: spectrograph/detector/data pipeline integration and
PFI/metrology integration.......................................................................................................... 64
Figure 5.1-2: Test Configuration of the Metrology, Positioner and PFI Loop at JPL...... 66
Figure 7.0-1: WFMOS Organization Chart............................................................................ 71
Figure 7.3-1: Top Level Schedule .......................................................................................... 78
Figure 7.4-1: Summary of Risk Assessment and Management........................................ 79
Table of Tables
Table 2.6-1: Dark Energy Surveys.......................................................................................... 16
Table 2.6-2: Galactic Archaeology Surveys .......................................................................... 17
Table 3.2-1: Requirements Traceability Matrix ..................................................................... 24
Table 3.4-1: Major Assumptions ............................................................................................. 26
Table 3.4-2: Assumed Telescope Performance Throughput Values ................................ 26
Table 3.5-1: Summary of Major Instrument Trades. The left column lists the trades
along with options considered. The selected option is shaded light blue. The right-handcolumn describes the rationale for the selected option........................................................ 27
Table 3.12-1: Spectrograph requirements based on the four surveys. The first threerows are a repeat of the science case. The other rows are derived requirements. ........ 47
Table 3.12-2: Basic derived spectrograph parameters. These are driven by the GAHR
survey, which is the most demanding of the four surveys................................................... 47
Table 4.1-1: Throughput budget for Dark Energy and Galactic Archaeology surveys.
The spectrograph efficiency differs for the two survey types because of differences in
the order sorting filter. The bottom line throughput number for the Galactic Archaeology
surveys reflects this difference. The throughput does not include optical coupling at thefiber entrance as this depends on seeing and the target (e.g. galaxy size). These factors
are accounted for separately in exposure time calculations. The entry for the WFC
vignetting is an average across the field. ............................................................................... 59
Table 4.2-1: Positioner Configuration Timeline .................................................................... 60
Table 4.4-1: Exposure and survey times for the low-redshift Dark Energy survey(DELZ), the high-redshift Dark Energy survey (DEHZ), the low-resolution Galactic
Archaeology survey (GALR), and the high-resolution Galactic Archaeology survey(GAHR). ....................................................................................................................................... 61
Table 7.2-1: WFMOS Costs by Partner ................................................................................. 73
Table 7.2-2: Cost by Top-Level WBS..................................................................................... 73v

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Table 7.2-3: Partner Shares and WFMOS Cost Distribution.............................................. 74
Table 7.2-4: Descope Options................................................................................................. 76
Table 7.4-1: Summary of Risk Assessment and Management.......................................... 80

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Executive Summary
The Wide Field Fiber-Fed Optical Multi-Object Spectrometer (WFMOS) will become thepremier tool for addressing two dominant themes in contemporary astrophysics: the nature of thedark energy (DE) responsible for the cosmic acceleration and the assembly history of galaxies
such as M31 and our Milky Way. Coordinated galaxy surveys delivering spectroscopic redshifts
for over 4 million galaxies will use baryonic acoustic oscillations and redshift-space distortions
to bring a major advance in our understanding of dark energy; the potential for the first breakthrough in this area is high. Equally ambitious Galactic Archaeology (GA) stellar surveys will
provide kinematic and chemical abundance data in order to identify and study the history andnature of galactic sub-components and their relation to dark matter halos surrounding galaxies.
This proposal is the outcome of our detailed WFMOS concept study. We have revisited andrebalanced the scope and design of the DE and GA surveys compared to the earlier Feasibility
Study. We have optimized the instrument for survey efficiency and flexibility. Our survey design
is complementary to the HyperSuprimeCam imaging survey, which will utilize weak gravitational lensing to constrain the growth of structure. With input from Japanese colleagues, we havealso examined a broader range of science applications ensuring our design is compatible with
their requirements.
In designing WFMOS, we have placed the greatest emphasis on survey efficiency and positioner reliability. Our prime-focus fiber positioning system comprises 2,400 actuated elementsthat fill a hexagonally-shaped 1.3 degree field of view. Piezo-electric motors rapidly translate
rigidly-mounted fiber optic tips without any tilting over 9.5 mm overlapping patrol regions,minimizing exposure overheads. Acquisition/guide and metrology cameras ensure accurate
placement and tracking with reference to a fixed set of fiducial fibers. High-throughput fiberconnectors convey signals to 3 identical spectrographs on the third floor of the Subaru Observatory. Each spectrograph is of a fast Schmidt design and permits a variety of moveable gratings
offering flexibility in spectral resolution and wavelength coverage. Thick deep-depletion silicon
CCDs provide enhanced red sensitivity for our proposed surveys.
Our team brings unique strengths to this ambitious project. The lead organization, JPL, isworld-recognized in the delivery of complex astrophysical facilities; it offers impressive technical depth in its engineering and management expertise. The Caltech-based PI has over two decades of experience managing observatories and instrumentation at the 4–10 m class telescope
scale. Our design team is drawn from a selection of institutes and individuals with excellent records in the design and delivery of innovative facility class instruments. Via the Thirty Meter
Telescope project, Caltech and JPL have established strategically-important links with the Subaru community, which ensure this proposal has strong institutional support. The science teamcomprises highly productive world leaders in observational cosmology and galactic structure,each with valuable experience in survey design and implementation. Excellent relations havebeen established with Japanese astronomers involved in faint object spectroscopy and the HyperSuprimeCam project.
Our cost for WFMOS is $58.3 M in real year dollars plus 20% reserves on the instrument fora total cost of $68.5 M. (These reserves are not applied to the Subaru provided elements and areseparate from the Gemini contingency.) This proposal focuses on accurate cost estimates, a clear
rationale for the level of reserves, and early identification of potential risks. The management
approach in delivering a complex instrument at cost with an international partnership draws onpractices developed over JPL’s history in the development and build of science instruments.

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WIDE FIELD FIBER-FED OPTICAL
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SUMMARY REPORTThis Study Summary encapsulates the results of our WFMOS concept study. We summarize
the development of our science objectives and requirements, present the instrument design features, trades and performance estimates, and discuss the implementation plan and management
and cost details for the construction phase.

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1.0

Introduction

The Wide Field Fiber-Fed Optical Multi-Object Spectrometer (WFMOS) is a fiber-fed spectrograph designed to be mounted at the prime focus of the Subaru telescope. It has two primaryscience objectives: 1) investigations of dark energy via baryonic acoustic oscillations and redshift-space distortions, and 2) the determination of the assembly history of the Milky Way andnearby galaxies through the study of stellar dynamics and chemical abundances (Galactic Archaeology). Such an instrument also has tremendous potential for discoveries over a muchbroader area of astronomy including studies of galaxy formation and large scale structure.
The WFMOS instrument is designed to allow simultaneous spectral investigations of 2,400
objects within a 1.3-degree field of view provided by the optical system of the Subaru Telescope.
The combination of a large primary aperture (8.2 m) and a rapidly reconfigurable multiplex system provides the capability to survey up to 100,000 objects in a single clear night. WFMOS will
use the Subaru primary mirror to focus light onto a focal plane incorporating 2,400 optical fibers.
These fibers relay light to three remotely located spectrographs (800 fibers per spectrograph),which simultaneously record 2,400 spectra on arrays of CCD detectors. Data are processed andarchived for further study and analysis by the WFMOS Data Reduction System.
The WFMOS instrument, described in Section 3, is comprised of the major elements shownin Figure 1.0-1, which serves as a useful introductory guide.

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WIDE FIELD FIBER-FED OPTICAL
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SUMMARY REPORT

Figure 1.0-1: WFMOS Instrument System Elements

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1.1

Proposal Organization
The Study Summary summarizes the concept study and includes an overview of the scientific motivation and survey strategy (Section 2), an instrument overview (Sections 3–5), operations description (Section 6), and the management and system engineering approach including
total cost (Section 7). The Management and Cost Volume provides project organization, costand schedule estimates, management techniques, and contract structure. The Detailed Technical
Design document is a detailed description of the instrument that supplements the instrument description in the Study Summary. The Operation Concept and Design Document presents detailed scientific context, science requirements, and operational mode and the Level 1 requirements. The Functional/Performance Requirements Document lists the instrument Level 2 andLevel 3 requirements.
Figure 1.1-1 shows an outline of the document tree with the current document shaded in blue.
Exec. Summary

Study
Summary

Mangmnt
and Cost
VolumeDetailed
Technical
DesignManagement.
Management proposal
including risk
management andsystem engineering
approach.

Technical. Full details
on each component ofour design study

Cost. Details of
proposal cost
including basis ofestimate andassumptions

OCDD

FPRD

Initial Operational
Concept Design
Document – Science
Context, Science
Requirements,
Operational Modes

Functional
Performance
Requirements
Document –Summary Instrument
Description,
Instrument Level
Requirements

Figure 1.1-1: Organization of submittal. Our concept study is summarized in the Study
Summary (this document). Separate volumes include: Management and Cost Volume, Detailed
Technical Design, the Operational Concept Design Document, and the Functional Performance
Requirements Document.

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2.0

Scientific Case and Requirements

2.1

Context
We can trace the motivation for an ambitious wide field, multi-fiber spectrograph such as
WFMOS to the early days of Gemini when the scientific promise of an f/6 Cassegrain focus on
Gemini-South was being enthusiastically promoted by various members of the UK, US, and Canadian communities. The f/6 secondary and its associated wide field corrector were eventuallyeliminated for financial reasons. However, the huge impact of galaxy surveys conducted using
multiplexed spectrographs on wide field 2.5–4 m telescopes such as SDSS and the AAT later
resurrected the need for an equivalent facility on an 8-m aperture telescope. The discovery ofdark energy in the late 1990s and the recognition that baryonic acoustic oscillations provide a
valuable tracer of the cosmic expansion, independently of supernovae, has given the issue further
impetus. The aim of this proposal is to make the strongest possible case for proceeding with theconstruction of WFMOS, in collaboration with our Japanese colleagues.
Our starting point in this study was the valuable Feasibility Study (FS) conducted by a consortium led by the Anglo-Australian Observatory. Gemini, Subaru, and the Gemini Board concluded such a facility should be installed at the prime focus of the Subaru telescope where it
could share the optics destined for the (then proposed, now funded) HyperSuprime Camera(HSC). The FS study proposed a range of ambitious surveys motivated by questions in both darkenergy and Galactic science. These questions, in turn, followed discussions by the Gemini community at the Aspen meeting in June 2003.
We have revised the science strategy proposed in the FS and, as a consequence, significantly
improved the design of the instrument. Recognizing that HSC is now funded, we have put considerable emphasis on harnessing the synergy between these two wide-field instruments. We
have also reassessed the role that WFMOS should play in dark energy studies. When the FS was
published, WFMOS was a rare example of a facility semi-dedicated to probing dark energy. Today, there are a number of potential probes using various methods with at least three (AAT WiggleZ, SDSS-3 BOSS and the Hobby Ebberley HETDEX) proposing a major spectroscopic effort
using baryonic acoustic oscillations. In the future there is optimism for one or two space missions dedicated to studying dark energy. Likewise, experience and progress with Keck’s DEIMOS spectrograph and the ESO VLT’s FLAMES spectrograph has significantly shaped ourthinking of the optimum parameters for Galactic science.
One thing has not changed since the FS. Subaru remains the unique platform for an instrument like WFMOS. Investigations by others have shown that such an ambitious wide field spectrograph could not be accommodated at Keck or the VLT. Short of building a dedicated widefield telescope, WFMOS on Subaru is the only route to achieving the science goals discussed inthis proposal.
The unique capabilities of the Subaru prime focus raise the important question of the attitudeof the Japanese community to WFMOS. Through active discussions and Japanese membershipof both WFMOS teams, this question has been brought into sharp focus over the past few years,
culminating in two major conferences—Panoramic Views of Galaxy Formation and Evolution(Hayama, December 2007) and Cosmology Near and Far—Science with WFMOS (Kona, May
2008). Recognizing the unique synergy between HSC and WFMOS on their national telescope,
the Japanese community views WFMOS as a general resource whose scientific impact should
extend well beyond that envisaged by the Aspen process or the FS. To address this, we co-opted
a group of 5 Japanese extragalactic astronomers enthusiastic about using WFMOS for programs
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not considered by the FS. We demonstrate in this proposal how our instrument, designed for darkenergy and Galactic surveys, is compliant with their needs.
The following sections summarize a more comprehensive discussion of the science goals, instrument requirements and proposed survey strategies contained in the Operational Concept Design Document (OCDD), to which the reader is referred.
2.2

Dark Energy and Modified Gravity
2.2.1 The Challenge of Cosmic Acceleration
The most radical development in observational cosmology in the late 20th century was thediscovery that the cosmic expansion has not decelerated as expected from the gravitational material it contains, but is now accelerating. The most direct evidence emerged from studies of distantSNe Ia used as standard candles. Additionally, by combining the now well-studied fluctuationsin the microwave background with measures of large scale structure, we can confirm independently that dark energy dominates the energy density of the Universe today.
To the extent that existing data can constrain the cosmic acceleration, it is consistent withthat expected for a non-zero cosmological constant Λ. However, adopting this explanation postpones solving the mystery, as no physical explanation can yet be found by physicists for such asmall value of the vacuum energy. The term ‘dark energy’ was coined to encapsulate our ignorance of the detailed physics being probed.
2.2.2 A Dark Energy Strategy
Dark energy can differ from a cosmological constant in being a dynamical phenomenon. To
test this we need to measure its equation of state through the parameter w = P / ρc2 where P is
the negative pressure and ρ its energy density. In principle w can also evolve: e.g., as a functionof the scale factor a, viz. w (a) = w0 + wa (1 - a). To constrain both the value of w (= -1 for acosmological constant) and to determine whether it evolves, requires precise measures of the distance - redshift D(z) relation in the large scale distribution of galaxies as a ‘standard ruler.’ High
precision is required: to measure w to 5% accuracy requires measuring the inferred distance D to1%.
An obvious strategy discussed widely in the community is to first undertake a concerted attempt to reject or otherwise constrain the cosmological constant (i.e., is w = -1.00?) leaving until
later the more challenging question of whether w evolves with cosmic time. A complication withmeasures based on the distance–redshift relation is that the bulk of the sensitivity to w arises notat the redshift at which the large-scale structure is being measured but, because of the integrated
signal that builds up from z = 0 where dark energy is at its most dominant, at an intermediate‘pivot redshift.’ The Kolb et al. Dark Energy Task Force advocated a figure of merit for a darkenergy experiment in terms of its ability to jointly constrain w(zpivot) and wa.
2.2.3 Baryonic Acoustic Oscillations
Baryonic acoustic oscillations (BAO) originate as horizon scales whose lengths are now precisely determined from the WMAP microwave background experiment. They appear in the
power spectrum of the galaxy distribution on scales of order 100 h-1 Mpc 1 . One of the principal
aims of WFMOS in the area of dark energy research will be to use the BAO signature as a standard ruler in order to measure w based on data at redshifts z ≈ 1 and beyond. This goal has been
discussed extensively in the literature, including specifically in the case of WFMOS, as well as
1

h refers to Hubble's constant measured in units of 100 km s-1 Mpc-1
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The fundamental limit to the precision with which we can measure the acoustic scale at any
redshift is sample variance. To be competitive it is necessary to survey large areas of the sky andprobe to faint limits. Both criteria favor WFMOS—it uniquely utilizes an 8-m–aperture telescope, and its 1.5 degree field of view dwarfs the competition (cf. 15 arcmin at Keck, 25 arcmin
for the VLT). Thus, WFMOS is uniquely capable of providing new constraints on fundamental
cosmology, particularly for redshifts beyond 1 where existing or planned wide field multiplexed
spectroscopic instruments (e.g., AAT, BOSS) cannot probe.
2.2.4 Modified Gravity and Redshift Space Distortions
Given the puzzlement about what dark energy might be, there is increasing interest in thecommunity in probing the validity of Einstein’s General Relativity. The degeneracy between
dark energy (e.g., a cosmological constant or a scalar field) and modified gravity is broken when
we consider a perturbed universe and examine the relationship between matter and the potentials.
In Einstein gravity the growth law for density fluctuations δ (= δρ/ρ ) is simple and can be represented via
fg(a) = dln δ / dln a ≈ Ωm(a)γ

where γ = 0.55. But for modified gravity, γ emerges as an important additional parameter in testing the standard framework.
The basis of our WFMOS dark energy proposal is two coupled galaxy redshift surveys designed to independently test the equation of state parameter w and the growth rate parameter γ.
We believe designing WFMOS and its surveys to measure the combination of both parametersoffers the greatest chance of an immediate breakthrough in our understanding of the cosmic acceleration. This is a major change from the strategy proposed in the FS.
We can test the growth of structure, and hence γ, using the clustering anisotropy on 10–20 h-1
Mpc scales measured using redshift space distortions (RSD). In redshift space, peculiar velocities
induced by perturbations add to the apparent Hubble motion and cause us to mislocate galaxies
in the radial direction. The 2-D power spectrum in redshift space is thus distorted from its form
in true space, an effect known as ‘Kaiser anisotropy.’ The effect was first seen with high signal/noise in the local 2dF survey using 220,000 galaxies, and the first tentative detections are
now being claimed at z = 0.8, opening the prospect for an initial test on γ (Figure 2.2-1).
RSD studies yield a measurement of fg at the redshift of interest. In contrast, BAO datameasure a combination of the local expansion rate H(z) plus the angle subtended by the acoustic
scale; this involves the angular diameter distance D(z), which is an integral function sensitive towhat happens to dark energy between redshift z and the present. But RSD measure peculiar velocities, which derive from the history of mass assembly at redshifts beyond that of the survey.
Thus, in combination, the BAO and RSD techniques test the continuity of whatever dark energy
might be over a wide range of cosmic time.

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Figure 2.2-1: (Left) The redshift-space correlation function ξ (rp, π) at <z> = 0.8, shown in false
color to emphasize the Kaiser anisotropy at amplitudes ξ = 0.1 from the recent survey of Guzzo
et al. (2008). The ‘finger of God’ radial stretching arises from virialized random motions at
separations of a few Mpc, indicating the degree of dark matter clustering at that epoch. (Right)
The growth rate function fg(z) inferred from redshift-space distortions using various existing andplanned surveys (black/blue points) and the precision that can be expected from a WFMOS
survey of 4 ×106 galaxies over 0.7 < z < 1.6 and 100,000 galaxies at z ~3 (red points, each split
into two redshift bins). See OCDD for discussion of the theoretical curves.

2.3

Proposed WFMOS Surveys
2.3.1 Survey Strategy: Changes Since the Feasibility Study
The FS proposed BAO surveys at both z ~1 and z ~3; the latter was an extremely expensive
endeavor given the longer exposure times. In contrast we have found that since the pivot redshiftchanges only modestly in between these two surveys (from z = 0.64 to 0.98), unless dark energy
is evolving rapidly, the case for an ambitious BAO survey at z ~3 is premature until we havesome indication of rapid evolution from the data at z ~1. On the other hand, since density fluctuations grow by a factor of 2 between z = 3 and 1, there is a good degree of independence between the growth rate information fg(z) at these redshifts. Consistent with our argument to constrain w and γ independently to the highest precision possible, we argue for investing most of theobserving time towards a 4 × 106 galaxy BAO survey at z ~1 that targets w (twice as ambitious
as that proposed in the FS) with a much more modest survey of 100,000 galaxies at 2.5 < z < 3.3(6 times smaller than proposed in the FS) primarily aimed at targeting RSD and hence γ.
Focusing on the z~1 survey, there is a further change since the FS, namely the funding of theSDSS-3 BOSS survey, which will define the landscape in low-redshift BAO studies over the
next few years. The BOSS plan is to invest a significant fraction of 5 years observing time toconstraining w using 106 luminous red sequence galaxies over z < 0.8. Although WFMOS as envisaged in the FS would offer a faster route to the same goal (or provide improved precision), we
argue that it is more effective to fully utilize the 8-m aperture of Subaru to complement BOSS by
tracing the distribution of a larger number of faint emission line galaxies to higher redshifts. By
adopting red-sensitive deep-depletion CCDs, we have extended the wavelength coverage of the
WFMOS spectrographs in order to follow the [O II] emission line in star-forming galaxies over
0.7 < z < 1.6. In fact, the improved efficiency of our spectrographs means that although we havedoubled the survey volume (cf. that proposed in the FS), the overall time requirement is compa-

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rable. In a survey spanning 4 × 106 galaxies, this component of our dark energy survey willmeasure the acoustic scale to about 0.7% precision in this redshift range.
Ahead of a potential JDEM/Euclid mission whose launch can be expected in ~2018–2020,
WFMOS will have no competition in the two proposed redshift ranges from any facility. The
only possible competitor is the HETDEX project, which plans to survey Lyman α emitters in ablind manner in the z > 3 regime. If that survey proceeds, our high-redshift survey would still bevaluable, complementing HETDEX by focusing on a different population of continuum-selectedsources.
In summary, the proposed combination of galaxy surveys discussed in this proposal will:• Measure the equation of state of dark energy using BAO over 0.7 < z < 1.6 to a precisionof 3% on constant w, better than the precision achieved by BOSS at lower redshift
z < 0.8.
• Measure the evolution parameter, wa, to a precision of approximately 0.3 (combining
with Planck and existing supernova data), offering an interesting test of possible
departures from a cosmological constant w = –1 over 0 < z < 2. This represents a DETF
figure of merit of close to 100, about 1.7 times better than expected from BOSS with thesame assumptions.
• Measure the growth rate fg(z) over 0.7 < z < 3.3 with unprecedented precision to test for
departures from Einstein gravity. The 0.7< z < 1.6 survey alone would measure fg toabout 1.5% precision offering the first convincing test of the standard model. The 2.5 < z< 3.3 survey would add a second constraint over a larger time baseline to 7%.
As discussed in detail in the OCDD, the 0.7 < z < 1.6 BAO/RSD survey will target 4×106
color-selected galaxies to RAB < 23.0 over a survey area of 3,200 deg2. Using the Keck DEEP2
survey data as a test case, we have developed techniques using BRI photometry that ensure ahigh success rate not only in selecting sources that lie in the required redshift range but also inindicating the sources likely to show strong [O II] emission, thus significantly improving thesurvey efficiency. Many ongoing and planned imaging surveys (DES, PanSTARRS, KIDS—in
which our team’s institutions are involved) will yield a photometric redshift accuracy better thanthat based on BRI photometry. The key issue is to ensure these surveys are sufficiently deep andlocated in an appropriate part of the sky for WFMOS with sufficient area. The simplest route is
to augment the HSC weak lensing survey (likely ~2,000 deg2) with DES, KIDS or Pan-STARRs
data. Additional HSC imaging will not likely be required (see OCDD for discussion). Spectroscopic completeness is not a concern provided the survey is uniformly selected over a large cosmic volume. With exposure times of 15 minutes and minimal reconfiguration overheads, thissurvey is expected to require 832 clear Subaru hours (~83 nights). 2
The 2.5 < z < 3.3 RSD survey will target 100,000 galaxies to RAB < 25 over a survey area of
100 deg2. Targets will be found using well-tested ugr Lyman break selection methods. gr (orequivalent) imaging data will be readily available from the HSC weak lensing survey; the morechallenging need is deep u-band data. Fortunately, the CFHT Legacy Survey has already obtained a significant (150 deg2) amount of u-band data with adequate depth. Some is already public, and no doubt the remainder will be in a few years’ time. With spectroscopic exposure timesof 4 hours, the high-redshift survey will require 286 clear Subaru hours (~29 clear nights).
Thus, in summary, the dark energy/modified gravity campaign can be accomplished in ap2 We adopt the Gemini-recommended unit that 1 night = 10 hours

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proximately 1120 clear Subaru hours (112 clear Subaru nights), an allocation that is surely
justified for such a profound cosmological quest. The expected achievements in constraining
both w (via the precision on D(z) ) and γ (via the precision on fg(z) ) with respect to existing andfuture facilities are summarized in Figure 2.3-1.

Figure 2.3-1: Fractional precision of various surveys. Left panel illustrates the expected
fractional precision of various surveys in measuring the acoustic scale (using our own S/Nformula, as discussed in the OCDD). Right panel shows the fractional precision in the growthrate (following White, Song & Percival 2008). The observed results for the 2dFGRS and SDSSat low redshift are plotted as points, together with the predicted accuracy of the VIPERS VLTsurvey. The lines are intended to show the precision expected in redshift shells of width 0.1. In
the left-hand panel, we also show (triangle points connected by dotted lines) the average
precision expected from the full volume of BOSS and of our WFMOS surveys.

2.3.2 Synergy with HyperSuprimeCam
HyperSuprimeCam (HSC) will provide an enlarged 1.5-degree diameter field and enter survey operations several years ahead of WFMOS. Our team has met with members of the HSC science community and discussed the obvious synergies between these two remarkable instruments.
Although scientists have yet to agree on the plans for using HSC, we use this proposal to stakeout two clear advantages of a coordinated effort.
Firstly, as discussed earlier, it makes great sense to coordinate the survey areas of both instruments and to use the HSC photometry to select targets for WFMOS. This involves very shallow imaging in 3 bands for some or all of the 3,200 deg2 z ~1 BAO/RSD survey where BRI (orgri) photometry has a proven track record in selecting 0.7 < z < 1.6 targets. Likewise for the z ~3
RSD survey, HSC data will be required to complement CFHT u data to select Lyman break‘dropouts.’
Secondly, the constraints on dark energy/modified gravity arising from WFMOS can, andshould, be combined with those determined using weak gravitational lensing constraints provided from HSC. Subaru/HSC will provide the best weak lensing data prior to the launch of
JDEM/Euclid. Combining weak lensing with photometric redshifts, calibrated via WFMOS spectroscopy, will enable 3-D weak lensing to constrain both the growth in the dark matter power
spectrum and cosmic geometry. We quantify the benefits of this remarkable combination in moredetail in the OCDD.
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2.4

Galactic Archaeology
2.4.1 Near-Field Cosmology Comes of Age
Understanding the formation and evolution of galaxies remains one of the key questions inastrophysics. The complex processes that govern growth over cosmic history are imprinted on
the Galactic structure see around us. WFMOS will be the premier instrument in the next decadecapable of connecting a detailed inventory of the Milky Way galaxy and its assembly history
with the rapid progress being made in probing the nature of the Universe on large scales and athigh redshift.
The questions posed by this area of ‘near-field cosmology’ are as fundamental as those posedin unraveling the mystery of dark energy. What can our Galaxy tell us about the validity of thestandard cosmological model and therefore the nature of dark matter and dark energy on small
scales? What role does dark matter play in the assembly history of galaxies and what processes
govern the interaction between dark matter and baryonic gas? Are these consistent with predictions of the standard model of cold, non-interacting dark matter? What can the fine structure of
the stellar distribution tell us about the distribution and nature of dark matter itself? What was thesequence and process of formation of the Galaxy and how does our understanding of the historyof the primary sub-components of the Milky Way (bulge, disk and halo and surviving satellites)
influence our picture of how normal galaxies evolved? Finally, from studies of neighboring systems, can we be sure our inferences from the Milky Way are applicable more generally?
The key to progress lies in developing a detailed picture of the fossil record of past events,
surveying stars to determine their spatial, kinematic and chemical properties. As galaxies are
dynamically and chemically inhomogeneous, only by using the higher dimensionality affordedby combining kinematic and chemical measures can a full understanding of their constituent
parts be determined. With its enormous multiplex advantage over an unprecedented wide field,
WFMOS offers unique opportunities in this area. As introduced in the Feasibility Study (FS),information of two kinds is required. A deep Low-Resolution (LR) survey will chart the structural composition of the Galaxy using kinematic information and metallicities derived from absorption line indices. A higher-resolution (HR) survey will use ‘chemical tagging’ to identify
sub-components as well as to unravel their detailed evolutionary history.
A number of developments have occurred in the subject since the FS was published in 2005.
We have taken full account of these, both in defining our science requirements for the instrument
and in fixing our proposed survey parameters. Foremost, we have placed significant emphasis on
synergies with the upcoming Gaia mission; this remarkable satellite will provide parallax andproper motion information to limits well-matched to those of our WFMOS surveys. The combination will deliver, for the first time, six-dimensional velocity phase space essential for identifying galactic components. Secondly, we have taken into account the early progress made in surveying nearby galaxies, both satellites of the Milky Way and the neighboring M31 and M33
systems. Much work is needed to extend these preliminary investigations that are crucial to determine whether the conclusions drawn from our Milky Way apply generally to most massive
spirals.
Finally, our team has been directly involved in progress made with Keck’s DEIMOS and
ESO’s FLAMES spectrographs. The latter, in particular, has demonstrated clearly the optimal
resolution for chemo-dynamical fingerprinting as a means to understanding the structure andevolution of nearby dwarf galaxies. As a result, our technical design and survey strategy havebeen significantly streamlined from those proposed in the FS.

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SUMMARY REPORTThe fundamental limit to progress in this area is the sheer number of high-quality spectra required. Utilizing Subaru’s 8.2-m aperture, WFMOS, with its the extensive simultaneous wavelength coverage, has unique leverage to address the key questions in Galactic Archaeology. (Figure 2.4-1 demonstrates very clearly that WFMOS simply has no competition in its informationgathering power.) 3

Figure 2.4-1: A measure of the impressive data-gathering power of WFMOS in comparison toother wide field spectroscopic facilities. The ordinate represents the product of the multiplex andnumber of resolved spectral elements above a fixed signal/noise per unit observing time. Our
team has learned much from its extensive use of the FLAMES and DEIMOS instruments, andthis has been reflected in our WFMOS requirements.

2.4.2 Survey Strategy: Differences with the Feasibility Study
The feasibility study provided an excellent starting point in highlighting the huge potential
gains available for Galactic science. However, the ambitious nature of the program led to a rather
imbalanced set of proposed surveys. We have reconsidered the program of Galactic Archaeology
in detail, balancing the scientific outcomes with a realistic instrument and survey design.
2.4.2.1

Low-Resolution Galactic Survey

For the Low-Resolution (LR) survey, synergy with Gaia is crucial since full-phase kinematic
measures and accurate distances are required for a complete analysis. Matching Gaia’s survey
limit of V = 20 is, therefore, a wise strategic choice. We have increased the spectral resolution
over the original FS for sound scientific reasons. Analysis shows that a resolution of R = 5000 isneeded to measure [Fe/H] to 0.1 dex precision and velocities accurate to 2 km sec-1; these are
required to provide an initial chemical tagging through overall metallicity measures and are es3

We define the rate of information as the product of the number of spectral resolution elements (R × λλrange), the number of fibers, the telescope diameter squared and the overall instrumental efficiency.

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sential to fully isolate and probe dynamically cold tidal streams and other substructures. Furthermore, spectra that simultaneously encompass both the Mgb complex and the Ca II triplet ensure the LR survey will provide valuable data for all types of star and will not only generate
more reliable stellar atmospheric analyses but also enable [α/Fe] abundance ratio estimates.
2.4.2.2

High-Resolution Galactic Survey

Experience with the ESO FLAMES and UVES instruments strongly suggests that the HighResolution (HR) aspect of the Galactic surveys should be undertaken at R = 20,000. This is optimal in terms of maximizing wavelength coverage, recognizing the intrinsic line broadening inmost stellar types. With a wide wavelength coverage and detailed line fitting, such spectra canrecover substantial abundance measures previously thought only to be accessible with higherresolution spectroscopy. At this resolution, our design provides a wavelength coverage (λλ4,800–6,800 Å) enabling detailed abundance analyses of a wider range of species and will also
facilitate improved links with spectral indices used in extragalactic studies.
In comparison with the original FS, where a HR survey of R = 40,000–50,000 was proposed,
our redesign results in consequent economies in spectrograph design and survey requirements
without significantly impacting scientific output. Furthermore, we dispensed with the ‘nod andshuffle’ technique proposed in the FS; as well as saving substantial CCD real-estate, our simulations and applications to real data demonstrate that a stable spectrograph with optimized sky subtraction leads to a factor 4 reduction in survey time.
The proposed WFMOS Galactic Archaeology program targets a number of Galactic andextragalactic components to address a range of scientifically important questions; the OCDD
(Section 4) provides further details of how our surveys are strategically partitioned across Galactic science ranging from halo streams, nearby dwarf galaxies, the central Bulge and the outerdisk and halo. In summary the LR survey totals 765 clear Subaru hours (~76 clear nights) andthe HR survey 1,125 clear Subaru hours (~112 clear nights). The section also presents the casefor future, large-scale surveys of M31 and M33; while this legacy program will require a substantial amount of further observing time, it will provide answers to the next big question in Galactic Archaeology, namely “Is the Milky Way unique?”In summary, our Galactic Archaeology surveys total 1,890 clear Subaru hours (189 clear
nights) and are more acceptable in their scope than those proposed in the FS, yet offer more detailed information. They are also more balanced with respect to the dark energy request.
2.5

Additional Science
In this section we discuss the additional science that will arise, either from the proposed surveys presented earlier, or via further observations.
2.5.1 Determining the Neutrino Mass
Neutrinos are so far the only dark matter candidates known to exist. Although their non-zeromass is firmly established from solar, atmospheric, reactor and accelerator neutrino experiments,
their absolute masses are still unknown. Through the suppression of the galaxy power spectrum
arising from neutrino free streaming, an upper limit to the sum of neutrino masses can be determined. The key point here is that the behavior of the matter spectrum P(k,z) with redshift z issensitive to neutrino mass in a manner that is distinctly different from that arising from the cold
dark matter component. The current upper limit, obtained using WMAP, the distant supernovae
data and BAO constraints from 2dFGRS and SDSS, is mv, tot < 0.61 eV (95% CL). WFMOS offers the prospect of reducing this upper limit significantly, possibly to the extent that the neutrino

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mass itself could be measured. As there is a degeneracy between the neutrino mass and the darkenergy equation of state parameter w, quantifying the neutrino mass is essential for most accurate
measurements of dark energy.
We have investigated the evolution of density perturbations in the presence of massive neutrinos and find that the nonlinear matter power spectrum at z ~1 can be particularly effective. The
nonlinear correction can improve the neutrino mass constraint by a factor of two compared withlinear theory (which underpredicts the suppression). Using our z ~1 BAO survey we forecast a 1sigma error of 0.097 eV on the mv,tot from WFMOS. This means that if mv, tot = 0.1 in nature,
WFMOS will be able to detect it at 1-sigma significance. This would be a development of profound importance for fundamental physics. On the other hand if the true neutrino mass is closer
to zero, then we would obtain an upper limit of 0.19 eV (95% CL).
Although an important source of uncertainty in constraining neutrino mass from galaxy surveys is biasing, fortunately there are several independent ways of controlling this systematic effect : (i) modeling the biasing via halo model or semi-analytic simulations, and marginalizing
over the biasing parameters; (ii) estimating the neutrino mass from galaxy power-spectra derived
for different galaxy types and checking for consistency; and (iii) estimating the biasing empirically from weak-lensing map to be produced (e.g., from HSC).
2.5.2 Studies in Galaxy Formation
A key conclusion arising from discussions with Japanese astronomers and presentations
made at both the international conferences mentioned in Section 2.1 is that WFMOS has enormous potential as a general-purpose, multi-object spectrograph in areas beyond dark energy/modified gravity and Galactic archaeology. One of the most exciting areas of great interest
to the Japanese community is the spectroscopic study of distant field and cluster galaxies, particularly those identified in the ambitious direct imaging surveys planned with HSC. Using theexisting Subaru imager Suprime-Cam, it has proved difficult for the Japanese community to secure necessary amounts of follow-up spectroscopy, primarily because the field of view of thecamera (30 arcmin) is so much larger than that of the multi-object spectrographs to which they
have routine access. The unique combination of HSC and WFMOS, instruments with nearly
identical fields of view, offers a new synergy in science areas where the Japanese community isalready very strong.
Recognizing this, we contacted an active group of Japanese extragalactic astronomers 4 andpresented our instrument design in order to see if it was compliant with their likely future programs. With their assistance, we developed a number of sample programs that are feasible with
our design. We hope this will serve to improve the reception to WFMOS in the Japanese community as well as to illustrate that survey instruments can achieve much more than their original
goals. Section 5 of the OCDD describes these programs in more detail.
2.6

Conclusion
In summary, our team has taken the case for WFMOS forward significantly from that proposed in the Feasibility Study.
The quest for dark energy has developed substantially since 2005. Observationally, many
new facilities have entered the race. Theoretically, the question of possible modifications to Einstein’s gravity is now being taken seriously. We have designed an instrument and two associatedsurveys that maximize the impact of WFMOS in both the equation of state of a possible scalar
4 M. Ouchi (convenor), T. Yamada, T. Kodama, M. Tanaka and K. Shimasaku

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field and in evaluating the validity of Einstein’s gravity via precision measures of the growth ofstructure. Our strategy takes maximum advantage of synergy with HyperSuprimeCam and complements the progress likely with the lower redshift survey undertaken with BOSS. Our surveys
will give valuable input into the likely programs that will eventually be conducted with JDEM.
The subject of near-field cosmology has progressed likewise. Using our experience withKeck’s DEIMOS and the VLT’s FLAMES, we have proposed a more balanced and economic
pair of Galactic surveys than those proposed in the FS. We have targeted the low-resolution survey closely to the upcoming Gaia mission, recognizing that synergy with parallaxes and propermotions will greatly enhance the science output. For the high-resolution survey, we have reduced
the spectral resolution, arguing that increased wavelength coverage will, overall, increase theyield. Dispensing with nod-and-shuffle detectors reduces the overall request to a practical number of Subaru nights.
The four surveys—Low-Redshift Dark Energy (DELZ), High-Redshift Dark Energy(DEHZ), Low-Resolution Galactic Archaeology (GALR), and High-Resolution Galactic Archaeology (GAHR)—are summarized in Tables 2.6-1 and 2.6-2.
Table 2.6-1: Dark Energy Surveys
Low-Redshift Dark Energy Survey (DELZ)
Redshift range
0.7 ≤ z ≤ 1.6
Number of galaxies> 4,000,000
Target density for successful redshifts
1,245 / deg2
Incorrect redshift fraction< 1%
Redshift precision, accuracyΔz < 0.001
Spectral range
6,300–9,700 ÅSpectral resolution
3,500
Magnitude limit
RAB < 23.0
Line S/N ratio
7
Estimated survey time
832 clear hours
High-Redshift Dark Energy Survey (DEHZ)
Redshift range
2.5 ≤ z ≤ 3.3
Number of galaxies> 100,000
Target density for successful redshifts
1,000 / deg2
Incorrect redshift fraction< 1%
Redshift precision, accuracyΔz < 0.001
Spectral range
4,200–6,500 ÅSpectral resolution
1,500
Magnitude limit
RAB < 25
Continuum S/N ratio
4
Estimated survey time
286 clear hours

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Table 2.6-2: Galactic Archaeology Surveys
Low-Resolution Galactic Archaeology (GALR) Survey
2,000 deg2
3 × 106
5,000
V = 20, S:N = 10-15 / Å
8,150–8,850 Å
4,800–5,500 Å
Wavelength centering
0.1 ÅContinuum estimation accuracy
5%
Relative flux calibration
5%
Estimated survey time
765 clear hours
High Resolution Galactic Archaeology (GAHR) Survey
Area
2,000 deg2
Number of stars
1 × 106
Resolution
20,000Sensitivity limit
V = 17, S:N = 150 / Å
Wavelength coverage
4,800–6,800 Å
Wavelength centering
0.1 ÅContinuum estimation accuracy
5%
Relative flux calibration
5%
Estimated survey time1125 clear hours
Area
Number of stars
Resolution
Sensitivity limit
Wavelength coverage

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3.0

Instrument Description

3.1

Rationale
The overriding factors in designing WFMOS for the science applications discussed in Section 2 are the need for survey speed and positioner reliability.
With the 8.2-m aperture of Subaru, emission line galaxies in the redshift range 0.7 < z < 1.6
can yield sufficient signal to noise for a reliable redshift in only 15 minutes. This places stringent
demands on the overhead for each exposure, and particularly that associated with reconfiguring2,400 fibers for a subsequent field.
In the course of conducting the surveys discussed earlier, WFMOS will reconfigure itself
over 3,000 times. Reliability of positioning and robustness of fiber units are key considerations.
In this section we present an instrument design that we believe represents a major improvement over that discussed in the Feasibility Study. The key advantages are:
• A positioner system that can reconfigure itself to the required precision in under 40 s.
Each fiber is maneuvered at its tip, eliminating the need for tilting. The stout design ofeach unit leads to a robust system.
• A set of three high-efficiency spectrographs based on a Schmidt optical design, which
offers an unrivalled combination of light beam area and acceptance solid angle at thedetector (the A-Omega product). Recognizing the long-term utility of WFMOS, ourspectrograph design can accommodate a range of gratings, each of which can be tilted,
offering full flexibility in survey design.
• State-of-the-art, deep-depletion CCD detectors, ensuring optimal efficiency at red
wavelengths for extending redshift surveys beyond z ~1.3 and measuring stellar dynamics
using the Ca II triplet.
3.2

Instrument Overview
WFMOS is a fiber-fed spectrograph designed to be mounted at the prime focus of the Subaru
telescope located on the summit of Mauna Kea, Hawaii. The instrument allows simultaneous
low-resolution spectral observations of up to 2,400 astronomical targets. In high-resolutionmode, WFMOS can observe 600 simultaneous targets.
The instrument system comprises several components. A block diagram appears in Figure
3.2-1. We summarize here the basic function of the components and their interactions; individualcomponents are described in greater detail in following sections.
WFMOS comprises components mounted in several locations. System elements required to
be at the prime focus of the telescope make up the Prime Focus Instrument (PFI) and are housed
within a Prime Focus Unit (PFU). The three spectrographs, along with their detectors and support systems, will be located in a spectrograph room adjacent to the telescope, above the Nasmyth focus.
WFMOS will use several Subaru provided elements that are being developed for the HyperSuprimeCam (HSC) instrument. These include the HSC Field Rotator, Hexapod, structure,
and a Wide Field Corrector (WFC).

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WIDE FIELD FIBER-FED OPTICAL
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Telescope Prime Focus Spider
Hexapod

Spectral Cal
SourceMetrology
CameraController

Metrology
CamerasInstrument Rotator
Spectral Cal Source
Controller

Spectrograph

Science Fiber
Illuminator

Actuator
Controller

Telescope Control S/W

Commands, Data & StatusSOSS

WFMOS Central Control
Computer
Control Software

STARS

OPE files(dataarchive)(observation
descriptions)

Optical fibers

Date Reduction

Electrical
ProvidedbySubaruFigure 3.2-1: WFMOS Block Diagram

19

Commands

Detector

Science Fiber
Illuminator
Data

Commands &StatusPFI

A&G
CameraController

Commands

SpectrographAcquisition & Guide
CamerasCommands

Fiber Positioners
Actuators

Detector
Commands & Data

Spectrograph

Fiber
Connector

Science Fibers

Detector

Science Fiber
Illuminator

Fiducial
Illuminator

Fiducial Fibers

Field Element

Wide Field Corrector

Alignment
SystemWIDE FIELD FIBER-FED OPTICAL
MULTI-OBJECT SPECTROMETER (WFMOS)

SUMMARY REPORTThe filter and Dewar window used with HSC are replaced by the WFMOS Field Element,
located close to the telescope prime focus. The Field Element is a flat window of 30 mm thickness and enables efficient coupling of the light from the WFC to the WFMOS fibers. It compensates for the removal of the Dewar window and filter that are part of the design for HSC. The
Field Element also contains an array of small obscuring spots under which unused spectrographfibers can be parked.
An array of 2,400 optical fibers is located in the PFU, with each fiber tip position being controllable in-plane by a piezo-electric Fiber Positioner, nicknamed “Cobra.” Each fiber tip can bepositioned within its 9.5-mm diameter circular patrol region. The patrol regions are in a hexagonal close-packed pattern, with 8-mm separation, and fill a hexagonally shaped 1.3-degree field ofview. The overlap between adjacent patrol regions enables 100% sky coverage of this hexagonal
field. The hexagonal shape of the field of view allows efficient tiling of the sky for large-areasurveys.
The fibers can be reconfigured for new observation fields in 40 s, with each fiber tip placed
to an accuracy of 5 µm (corresponding to less than 0.054 arc seconds on the sky). Fiber tips aretranslated (rather than tilted) within the image plane, providing uniform coupling efficiency for
all configurations.
The small fraction of fibers that is not allocated to astronomical targets will measure thespectrum of the sky. In the case of the GAHR survey, which is cross-dispersed, three out of
every four fibers are obscured by positioning the fiber tip behind a small dot on the back of the
field element to avoid overlaying spectra or adding unnecessary sky background.
The Fiber System consists of 2,400 fibers that relay light incident at the instrument’s focal
plane to three spectrographs located remotely in a room adjacent to the telescope. The fibers are
60 m in length and are divided into three groups, each providing 800 inputs to a spectrograph. A
key component of the fiber system is a Fiber Connector, which allows most of the fiber bundle
to be permanently installed to the telescope while facilitating the removal of the PFU as necessary. The connector is located 5 m from the positioner end, allowing removal of the Top End
Unit while leaving 55 m of the fiber bundle on the telescope structure.
After the science fibers enter the spectrograph room, they are divided three ways and arranged as linear arrays of 800 fibers at the input to each of three identical Spectrographs. The
spectrographs provide the capability to simultaneously measure 2,400 spectra in low-resolution
mode and 600 spectra in high-resolution mode. During high-resolution spectral observations, theunused fibers are positioned underneath Blocking Spots to reduce the sky background contribution to the spectra in the high-resolution mode. Each spectrograph uses a Schmidt optics systemcomposed of a collimator, grating, and camera optics that are reconfigurable to support the four
primary science surveys, which require different spectral resolutions and spectral ranges.
At the image plane of each spectrograph is a Detector System consisting of a pair of close
butted 2,880 × 5,760 pixel CCDs inside a Dewar and associated cooling system. The 800 spectra
are dispersed along the columns (5,760 pixels), so the gap falls between spectra. In the case of
the GAHR survey, the 200 cross-dispersed spectra fill the detector array. The CCDs have exceptionally high quantum efficiency, which improves overall instrument efficiency.
The Acquisition and Guide (A&G) System is composed of four cameras located on the instrument optics bench in the PFU. Once calibrated to the fiber optic focal plane, the Acquisition
and Guide System will provide feedback to update the telescope’s pointing and field rotation bytracking on guide stars within the cameras’ field of view.

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There are four Metrology cameras, mounted on the prime focus support struts looking down
at back-illuminated fiber tips via the primary mirror of Subaru. The metrology cameras determine the location of the fiber tips, allowing accurate positioning of the fibers on the selected science targets. Each metrology camera consists of imaging optics with a CCD detector and includes an internal flat-field source. Calibration of geometrical distortion inherent in the optics ofthe metrology cameras is achieved by having 154 illuminated fixed Fiducial Fibers whose positions are known accurately. Of the fixed fiducials, 91 are within the bounds of the hexagonal
field of view, while 48 fixed fiducial fibers are at the perimeter of the hexagonal field. Beyond
the perimeter of the hexagonal field, 15 fiducials are used to provide encoding of the rotator angle. The fiducial fibers also allow correction for geometric distortion in the WFC, through measurement and modeling.
The field of view of each Metrology camera is designed to image 1/4 of the optics bench fibers. The cameras image the backlit science fibers by viewing the optics bench through the telescope optical system, including the WFC and the primary mirror. The centroid of each fiber image is calculated and used to determine each science fiber position with respect to the fixed-fiber
locations. This coordinate system is accurately referenced to the A&G camera, allowing accurate
placement of astronomical targets on the science fibers.
The science fibers are back-illuminated by sources located inside the three spectrographs.
Only one set of 800 fibers is illuminated at a time, in a manner allowing no ambiguity in determining fiber positions in the case that two fibers occupy a region of overlapping patrol regions.
The Prime Focus Instrument (PFI) consists of the Optical Bench, which contains the Field
Element, Fiber Positioner System, Acquisition and Guide System, and components of the Fiber
System Power Distribution System and Thermal System. The optical bench is mounted to the
Rotator within the PFI via a bi-conic structure. This structure is tuned in stiffness to compensate
for differential flexure between the WFC and the Rotator interface where the WFMOS components are mounted. The PFI is housed within the PFU, which can be installed and removed fromthe telescope as a single unit.
The Power Distribution System is located within the PFU to provide power conditioning
and control to the electronics associated with the Fiber Positioner System, the Acquisition and
Guide System, and the Metrology System. The distribution system interfaces to the telescope
facility power and receives power control inputs from the System Control function. The Thermal System provides the interface between the telescope facility cooling system and the instrument systems located in the PFU to prevent excess heat from disturbing the telescope’s viewing
environment.
The Spectral Calibration Lamps provide the capability to measure the spectral response of
the WFMOS instrument by providing a spectral source at the input to the telescope. The calibration lamps will be mounted on the telescope apart from the Prime Focus Unit.
The System Control provides the capability to manage an observation’s preparation andexecution. Preparation tasks include parsing large surveys into a series of individual observations, matching fibers to designated targets and selecting appropriate guide stars, predicting exposure times, and describing observations in a Subaru-format file. During telescope observations,
the System Control orchestrates the instrument systems to achieve successful data collection, sequences the Metrology and Fiber Positioner Systems to achieve target acquisition with the fibers,interfaces to the Subaru Observatory Software System, and provides operator interfaces forquick-look data and instrument health and safety displays. The System Control also provides the21

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capability to perform high-level engineering tests to aid in instrument integration and troubleshooting.
The Data Reduction function receives raw data read from the Detector System and uses it toproduce calibrated spectra ready for scientific analysis. The Data Reduction function provides
the capability to reduce data of varying spectral resolution and saves all spectroscopic and image
files in the required data format. The Data Reduction function will also provide the capability toquickly assess the quality (completeness, signal-to noise ratio, etc.) of the data collected.
A typical observation proceeds as follows. After initialization, instructions are executed toconfigure the telescope and instrument for the desired observation. The details of each observation are pre-planned using software provided by the System Control to aid the researcher in orchestrating the survey. Using information derived from a catalog of astronomical objects, thefield configuration, including the positions for each of the 2,400 fibers, is computed prior to thestart of the observation. The spectrograph configuration is determined by which of the four primary science surveys is being executed. A typical observing night might use more than one spectrograph configuration in order to optimize survey efficiency as sky brightness or the seeingchanges.
Execution of an observation begins with moving the telescope to the desired field center.
While the telescope is slewing, the Fiber Positioning System executes commands from the System Control to simultaneously move each of the 2,400 fibers to its required position. The Metrology System is used to verify each fiber’s position. The positioning process proceeds as follows:
1. The current position of each science fiber is determined relative to a set of fixed fiducial
fibers. These fixed fibers are backlit by a dedicated calibration lamp. The backlit fibers
are imaged by the Metrology System and the centroid of each fiber is calculated. Similarly, the 2,400 science fibers are backlit by calibration lamps in each spectrograph andimaged by the Metrology System. The centroids of each science fiber image are computed in order to determine their relative positions.
2. The System Control computes and issues the necessary commands to the Fiber Positioning System to move the 2,400 science fibers to the positions required for the observation.
The metrology and the fiber movements are coordinated in 3 groups of 800 fibers. Some
operations proceed in parallel.
3. Steps 1 and 2 are repeated several times to accurately position each fiber.
The positioning process is estimated to take less than 40 s from start of telescope pointing tocompletion of fiber positioning. In parallel with this process, the Acquisition and Guide System
refines the pointing of the telescope. The Acquisition and Guide System provides feedback to thetelescope pointing and rotator position during the observation. The hexapod is commanded to theappropriate position to maintain the WFC alignment in the presence of flexure in the PFU.
Once positioning is complete, the Detector System collects data. The Data Reduction Systemproduces quick-look information on the completed observation so to that survey progress can bemonitored.
3.3

Major Requirements
The OCDD defines the science requirements for WFMOS. These requirements in turn define
the instrument requirements, as described in the FPRD. The WFMOS FPRD has over 60 requirements on the entire instrument and several hundred requirements on the subsystems. Not all

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requirements are of equal weight; some impact the design much more than others.
We have identified those requirements that govern the design and established their connection with the WFMOS science requirements in Table 3.2-1, the Requirements Traceability Matrix. The table traces the flowdown of science objectives to measurement requirements, and thento instrument properties. The final column lists the instrument expected performance.
These requirements are abbreviated versions of the full requirements listed in the OCDD and
FPRD, and are intended to summarize succinctly how high-level science requirements drive thedesign, and to identify the subsystems affected by high-level requirements.

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Table 3.2-1: Requirements Traceability Matrix
Scientific Measurement
Requirements

Science Objectives
Provide a constraint on w to 3%.

Measure the distance to amean redshift 1.2 with anaccuracy of 0.6%.

Instrument Functional
Requirements

Predicted Performance• Redshift of the OII doublet at 373 nm
from a redshift of 0.7–1.6• Optimize visibility of OII doublet in OH
night sky spectrum• λ = 630–970 nm• R = 3500
• Line S/N ≥ 7• λ = 630–970 nm• R = 3500
• Line S/N = 7• Observe > 4 × 106 galaxies• ≥ 2,400 fibers• 2,400 fibers

Less than 2000 hours for
DELZ and DELHZ• Fast observations and high efficiency• 80 s or less between
observations• ≥ 2,400 fibers• 40 s between
observations• 2,400 fibers

Measure the growth rate of density
fluctuations.

Measure this in 4
independent redshift bins
between redshifts 1–3 to a
precision of 0.02 nearredshift 1 and a precision of
0.1 near redshift 3.

• > 100,000 galaxies observed• < 1% incorrect redshift• Redshift precision: dz < 0.001• λ = 420–650 nm• R = 1500 ± 5%
• Continuum S/N ≥ 4• λ = 420–670 nm• R = 1500
• Continuum S/N = 4

Measure the mass, extent,
substructure, of the Milky Way dark
matter halo

Characterize radial velocity
in approximately 30% of
substantial stellar streams in
the halo.

•
•
•
•
•

Observe > 3 × 106 stars w/ V ≤ 20
R ~5000
λ = 480–550 nm and 825–885 nm
Wavelength centering ≥ 0.01 nm
Continuum estimation ≥ 5%

• ≥ 2,400 fibers• R = 5000 ± 5%
• λ = 480–554 nm and 818–885nm• Spectral calibration lamps emit
over λ = 480–544 and 818–885
nm.
• Scattered light < 5%

• 2,400 fibers• R = 5000
• λ = 480–554 nm and
818–885 nm• Spectral calibration
Lamps emit over λ =480–544 and 818–885
nm.
• Scattered light < 5%

•
•
•
•
•

Observe > 1 × 106 stars w/ V ≤ 17
R ~20,000
λ = 480–680 nm
Wavelength centering ≥ 0.01 nm
Continuum estimation ≥ 5%

• ≥ 2,400 fibers• R = 20,000 ± 5%
• λ = 480–680 nm• Spectral calibration lamps:
λ = 480–544 and 818–885 nm.
• Scattered light < 5%

• 2,400 fibers• R = 20,000• λ = 480–680 nm• Spectral cal lamps:
λ = 480–544 & 818–885 nm.
Scattered light < 5%

• 80s or less between
observations• ≥ 2,400 fibers• 40 s between
observations• 2,400 fibers

Conduct kinematic measurements ofdSph candidates to confirm status, &measure masses, density profiles,
and extent of dSphs.
Measure the extent, (sub)structure,and kinematics of Milky Way stellar
populations
Explore the fossil record of chemical
evolution of stellar populations in the
Milky Way.

Radial velocity precision of 2
km/s
Proxy [Fe/H] measurements
to 0.1 dex

Measure metalicity of stars
with equivalent widths tobetter than 5 mÅ

• Fast observations and high efficiency
Less than 2500 hours for
Galactic Archaeology Survey

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3.4

High-Level Assumptions
This section describes the high-level assumptions used to finalize the proposed design of
WFMOS (Table 3.4-1). A full list of assumptions appears in the Detailed Technical Design.
We received documents defining the performance of the Subaru Wide Field Corrector(WFC); however, we do not have design details that can be used in the system-level analysis.
Based on the conceptual design provided, which shows the number of optical elements andqualitative curvature of each optical surface, we developed a working model of the WFC consistent with the provided performance specifications. The mechanical tolerances in our positioner
design and the design of the field element are based on our assumed WFC model.
We assumed that a portion of the PFU would be shared between HSC and WFMOS, that the
instruments would need to be de-mated from the PFU before the other instrument can be mated,
and that this process would take approximately a day. If it takes substantially longer than a day,it will affect the observing timelines. We also assumed that it would only take three people toremove or install WFMOS from the telescope and move it to the TUE room. We further assumed
that Subaru could staff this task without hardship. A related assumption is that we have two interfaces between WFMOS and the PFU: one at the rotator and one at the top of the hexapod.
Another important assumption concerns the stiffness and mechanical design of the HyperSuprimeCam PFU. We assumed that the instrument hexapod can use a static look-up table to compensate for the flexure in the PFU structure. We further assumed that there will be non-negligible
differential flexure between the WFC mounting interface and the interface where the fiber positioning system must mount. Further details are given in Section 3.9.
We assumed the performance parameters for some aspects of the Subaru telescope and theelements common to HSC and WFMOS would not dominate the error budgets and assumed values were used as shown in Table 3.4-2. As a measure of caution, whenever we encountered an
unknown parameter, we increased the level of margin held in the budget.
We assumed that the HSC wavefront sensor will produce lookup tables that will control the
PHSC hexapod accurately enough for WFMOS, as suggested in the Interface Control Documentation (90.00.00.01_15.10.00.00_ICD_PHSC). Our computations using the reversed engineered
WFC model showed that we need the focus of the WFC to be accurate to 50 μm.
We assumed that we will be able to send focus updates to the hexapod to adjust the focus of
the WFC relative to the Subaru primary. We likewise assumed we will be able to send updates
and offsets to the Instrument Rotator with a high speed communications link. Our Acquisition &Guide System will modify both focus and rotation.

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Table 3.4-1: Major Assumptions
Assumption
Differential flexure between WFC
interface and positioner interface
WFC vignetting and telecentricity

Basis
HSC design and WFC mass

PFU interface
PHSC hexapod performance andstability
Ability to send focus updates to thehexapod to adjust the focus of the
WFC relative to the Subaru primary.
Ability to send updates and offsets tothe instrument Rotator

Technical Impact
Need for a passive flexure compensation
mechanism.
Field corrector design

WFC performance provided, JPL
model of WFC
HSC design
HSC design

Fiber placement accuracy
Fiber placement accuracy and stability

HSC design

Optical coupling efficiency to fibers

HSC design

Fiber placement accuracy and stability

Table 3.4-2: Assumed Telescope Performance Throughput Values
Component /Subcomponent
Atmosphere
Secondary Obscuration
Primary Reflectance
WFC Throughput
WFC Vignetting
Uncorrected Atmospheric
Dispersion
Misalignment between
WFMOS/HPFU

420 nmComponent
Throughput
0.83
0.95
0.87
0.88
0.83

440 nmComponent
Throughput
0.83
0.95
0.87
0.88
0.83

550 nmComponent
Throughput
0.89
0.95
0.87
0.88
0.83

640 nmComponent
Throughput
0.90
0.95
0.85
0.88
0.83

790 nmComponent
Throughput
0.90
0.95
0.82
0.88
0.83

970 nmComponent
Throughput
0.90
0.95
0.82
0.88
0.83

0.995

0.995

0.995

0.995

0.995

0.995

0.95

0.95

0.95

0.95

0.95

0.95

We assumed we will get the following documents at the start of any contract. These assumptions have driven our cost and schedule estimates for the build phase.
• Complete Subaru Interface Control Documents.
• Optical Prescription to the Wide Field Corrector.
• Complete Subaru and Gemini requirements.
• Finite Element Analysis for the WFC/PHSC interface.
• Dynamic environment from telescope movement at spectrograph room and at Prime
Focus
During the course of the build phase we assumed we will get:
• SOSS Software Interfaces by Control System PDR.
• The Gemini Pipeline Reduction System CDR Documentation by Data Reduction Pipeline
PDR.
• The Gemini Pipeline Reduction System Software by Data Reduction Pipeline CDR.

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3.5

Instrument Trades
An important aspect of any study is tracing the high-level design trades that lead to the proposed design. Table 3.5-1 provides a summary of the trades considered for various subsystems
with the selected option highlighted and a brief description of the rationale for that choice.
Throughout our design, we have optimized throughput subject to cost and risk constraints.
Table 3.5-1: Summary of Major Instrument Trades. The left column lists the trades along with
options considered. The selected option is shaded light blue. The right-hand column describes
the rationale for the selected option.
Rationale

Trades and Selected Option
Fiber Size
107 μm core (custom)
100 or 110 μm core (standard)
Detectors
Custom 450 μm LBL w/ delta doping, custom
TiO2/SiO2 coating
300 μm thick Hamamatsu CCDs
Spectrograph Design
Reconfigurable dichroic system
Schmidt configuration - VPHG reconfigurable,
identical
Positioner
Articulate fiber tip—Cobra
Articulate fiber base
Robot pick and place
Fixed fibers, actuated mirror to direct light into
fiber
Plug plate
Metrology Cameras
Four 4.5k by 4.5k back illuminated full frame
CCD detectors
Four 4.5k by 4.5k front illuminated full frame
CCD detectors
Two 2k by 2k back illuminated full frame
detector scanned by several hundred pixels
Single large format 10.5k by 10.5k pixel backilluminated full frame detector
Corrector Plate
Plano-plano
Spherical

Optimizes fiber performance with low (3%) cost increase over standard fibers.
Numerical aperture is 0.22

Increased performance w/ custom CCD. Estimated QE > 90% between 450970nm, > 96% between 480–950nm

Major advantage is reduction of complexity and number of optics. Large beam
size minimizes effect of blockage and allows increase in resolution if desired.
Places control on the point of interest -fiber tip. Does not require
corresponding optics (POSM)

Low risk, cost effective
Perform early testing. Cost savings over back illuminated version if
requirements are met

Simple, inexpensive, equivalent performance

3.6

Metrology Camera System
The function of the metrology system is to measure the position of the fiber heads on the focal plane so that each can be repositioned to a given location to an overall accuracy of 10 μm(L2.300). The metrology system itself must be able to locate the fibers on the positioner focal
plane to better than 3.2 μm (L3.2720).
The proposed system has four cameras, mounted on the prime focus spiders looking back at
the back-illuminated fiber heads via the primary mirror of Subaru (Figure 3.6-1). Each camera
27

WIDE FIELD FIBER-FED OPTICAL
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SUMMARY REPORT

has a 4k × 4k CCD detector with imaging optics and will look at a quarter of the positioner focal
plane. Calibration of the camera distortion and mapping to the focal plane will be achieved byimaging backlit fixed fiducial fibers of known position located on the prime focus positioner
plane.

Figure 3.6-1: Schematic of Single Camera on the Prime Focus Support Strut

To reposition the science fibers onto a new object field, each metrology camera takes an image of a quarter of the focal plane. This is achieved by pulsing the LED light sources of the backlit science and fiducial fibers.
Since the patrol fields overlap, multiple images of the focal plane with different illuminatedsets must be taken in order to identify the fibers uniquely. One-third of the science fibers arebacklit at a time using an LED source in the spectrograph, and metrology imaging is performed
three times. The science fibers are then moved to new positions based on predicted positions of
the chosen targets. The metrology camera then takes a new set of images, and the positions of thescience fibers are measured in the same way as before. The error from the required position isthen calculated and the fibers moved again. This process is repeated six times.
The stringent subpixel centroiding requirement means back-illuminated CCDs are preferred
to front illuminated ones, as they offer better subpixel uniformity. The baseline detector is the
1100S camera from Spectral Instruments, Inc with a E2V230-84. The camera read-out time is 2s, and the centroid computation time is 1 s.

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WIDE FIELD FIBER-FED OPTICAL
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Figure 3.6-2: 1100 Series Camera from Spectral Instruments, Inc.

The optical design of the camera and its position on Subaru is shown in Figure 3.6-3. Since
the camera is looking back at the fiber field via the primary, the camera is focused at infinity.
The system is designed to be compact while giving very good image quality across the field withlittle field distortion. The CCD detector window is coated with a narrow-band filter to reduce
chromatic and stray lights effects. The backlit fiber sources are matched to the narrow band filter
and are visible for daytime work on the telescope.
TelescopePrimary
Wide Field
Corrector

Metrology
CameraFigure 3.6-3: Optical Layout of the Subaru Telescope and the Metrology Camera

3.7

Field Element

The Wide Field Corrector design assumes a filter and Dewar glass window of 30-mm diameter. When WFMOS is installed, the WFC remains, but the filter and Dewar window inherent in
HSC are removed. A trade study was performed to consider whether alternate designs would
provide significant advantage over the flat window. The factors considered were flat image
plane, coupling efficiency, ease of fabrication, image quality, tolerances and ease of applying
blocking spots. While no replacement window would be the least expensive, a simple shift in the
focal plane is not sufficient to provide good coupling to the fibers, due to the residual sphericalaberration and field curvature. Our design, therefore, incorporates a flat optical element thatserves as a replacement for HSC’s Dewar window and filter.
An important aspect of the optical performance of the WFC is the lack of “telecentricity”: thechief ray is not normal to the focal surface at the edge of the field. We evaluated the loss of fiber

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WIDE FIELD FIBER-FED OPTICAL
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coupling efficiency due to this effect and the potential of additional optical elements to compensate. Through our optical modeling, we have determined that the loss of throughput due to lackof telecentricity is very nearly identical to the loss of light from vignetting in the WFC. Additional optical elements designed to improve telecentricity do not actually significantly improve
the throughput and have, therefore, not been incorporated in our design.
The WFMOS Field Element coupling efficiency at the fiber tips is shown in Figure 3.7-1.
The blue line indicates the coupling efficiency without the as-built performance of the Subaru
Primary and WFC. The red line gives the overall coupling performance including the Primary,
WFC and average seeing of 0.7 arc seconds. The black line shows the goal of 0.80 coupling efficiency.
In addition, the Field Element has 2400 small, opaque spots (“Blocking Spots”) deposited on
the side closest to the fiber positioners. There is one Blocking Spot per patrol region and each
Blocking Spot obscures a tiny fraction (<0.1%) of each positioner patrol region. The presence of
the Blocking Spots allows the positioner to park unused fibers underneath these spots duringhigh resolution spectral observations, which in turn reduces the sky background during these observations.

Figure 3.7-1: The WFMOS Field Element–to–Fiber Tip Coupling Efficiency

3.8

Acquisition and Guide Camera
The WFMOS Acquisition and Guide (A&G) System serves two purposes: during initial target acquisition, it locates the appropriate guide stars; once the stars have been acquired, the A&G
system sends positional updates to the telescope and the image rotator in order to maintainalignment of the fibers with the science target.
We have based the A&G system for WFMOS on the highly successful A&G system for
Autofib-2 on the William Herschel Telescope on the Canary Islands (Goodsell et al. 2003; Parryet al. 1994). WFMOS will use four coherent fiber bundles, each made up of thousands of individual fibers. Each fiber bundle relays the image from the focal plane to one of four CCD cameras. The image of the fiber bundle needs to be reimaged onto the detector. Figure 3.8-1 summarizes the configuration of the A&G system for WFMOS.

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WIDE FIELD FIBER-FED OPTICAL
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Figure 3.8-1: The proposed configuration of the WFMOS A&G System. A coherent fiber bundle
reimages the focal plane onto the A&G Camera.

3.9

Prime Focus Instrument
The elements just described—the Acquisition and Guide Cameras, the field element, thepositioner subsystem and the other telescope-mounted components except the metrology cameras—form the Prime Focus Instrument (PFI) and are mounted on the PFI structure that interfaces to the PFU) (Figure 3.9-1). The critical function of the PFI structure is to establish andmaintain alignment of these subsystems to the WFC. There are two major elements of the PFI
structure: the rotating portion that interfaces to the PHSC rotator and the fixed portion that interfaces near the upper end of the PHSC hexapods.

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Figure 3.9-1: The Prime Focus Instrument

The rotating portion interfaces to the PHSC Rotator with a bi-conic structure and supports the
Positioner Assembly and associated fibers, the fixed fiducial system, and the acquisition andguide cameras (Figure 3.9-2). The bi-conic structure provides initial alignment of the positioner
to the WFC and is tuned in stiffness such that the defocus, decenter, and pointing of the Positioner Assembly matches that of the WFC within the required limits. The most critical alignment
requirement is the focal plane tilt of 34 arcsec with respect to the WFC. This appears to be readily achievable since this is on the same order as the predicted WFC tilt with respect to the telescope over its 60 degree operational elevation.

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Figure 3.9-2: The Rotating Portion of the PFI

The PFI Fixed Plate supports the optic fiber cables, the Strain Relief Box, power and data cables, and the Cable Wrap Assembly (Figure 3.9-3).

Figure 3.9-3: The PFI Fixed Structure

The Fiber Optic cable is routed between the fixed and rotating parts of the PFI, such that the
500-degree rotator range is accommodated by twisting along the rotation axis rather than bending around a spool. This improves throughput by reducing the bundle length by approximately 4m, while the effect on throughput and FRD via twisting is offset by the cable not having to bendaround a spool or cable chain.
The thermal control approach is to remove heat locally from the heat-generating componentsof the system—the A&G Cameras, and the Metrology Cameras—using the glycol coolant loop.
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These components are commercially available with glycol cooling, which may be modified toremotely locate the field joint on the secondary ring.
One of the major integration phases occurs at JPL where the positioner assembly, Acquisition
and Guide Cameras, and Fixed Fiducial are integrated onto the PFI structure using a temporary
structure for the PHSC. As a risk-mitigation strategy, a fit check of the Rotator Interface Ring
and the PHSC Rotor is done before it is integrated into the instrument. A turnover/rotation fixture is used to build up the PFI starting from the rotator interface. As soon as the alignment system and positioners are integrated onto the Rotator Interface Ring, the fixture is used to test performance in multiple orientations with respect to gravity.
3.10 Positioner
One of the key features of WFMOS is the positioner assembly, which places the tips of thefiber optics at specified locations on the focal plane to route the light from the astronomical objects of interest to the spectrograph. Previous generations of multi-object spectrographs have allowed simultaneous spectra of many hundreds of objects. Techniques employed for these spectrographs have included humans placing fibers by hand in automatically machined plates (e.g.,
SDSS), and robotic pick-and-place of magnetically attached fibers (e.g., 2dF). As our ambitions
rise to simultaneously measuring many thousands of simultaneous spectra, a reconsideration of
the full array of viable techniques is required.
After examining a number of alternatives (Detailed Technical Design, Section 2.3.16), the selected approach, the “Cobra” positioner system, populates the focal plane with 2,400 positioners
that fill a hexagonal field (Figure 3.10-1). Each positioner moves a single fiber to its target position in its planar patrol region. Translating the fibers through the patrol region using these actuators allows rapid, direct placement of the fiber tip with a mechanism that is robust in terms ofstability, accuracy, and life. Its planar movement eliminates losses due to fiber tilt and tests replicating the fiber path through the positioner confirm that there is no loss due to fiber twist or
bend.

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Figure 3.10-1. Positioner Bench Populated with One Module and a Single Positioner

Each fiber tip can be placed anywhere within a circular region, the patrol region, using a setof two piezo-electric motors. Figure 3.10-2 shows how the focal plane is covered with overlapping patrol regions. In this figure, the blue segment and the circle it sweeps out shows the travelof the first (theta) stage; the orange segment and its swept circle shows the coverage of the second (phi) stage. The second stage rotates the fiber arm. The combination of these movements
provides coverage through the larger enclosed circle with radius, R.

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Figure 3.10-2: Positioner Patrol Region Overlap. The black dashed lines indicate the patrol
regions of adjacent positioners, showing there is complete coverage of the populated field ofview. The distance between adjacent positioners, d, is 8 mm, while the diameter of the Patrol
Region, 2R, is 9.5 mm. Blue and orange represent the two stages of the positioner motors (seetext).

The positioner is designed to locate the fiber tip to within 5 μm accuracy in less than
10iterations (FPRD L3.790), with all commanded moves taking less than 1 second (FPRD
L3.800). Meeting these timing and accuracy budget allocations means the throughput loss fromthe positioners is less than 10%. Fiber repositioning — taking into account the other steps suchas the metrology loop (Table 4.2-1) — will be accomplished in less than 40 seconds. This maximizes observing time and efficiency. A prototype was built by New Scale Technologies and recent testing at JPL demonstrated this first unit was able to converge on 90% of the targets towithin the 5 microns accuracy in 6 iterations and had 100% convergence after 8 iterations, all at
the required speed. Additional details of this test are discussed later in this section and in the Detailed Technical Design, Section 2.3.14.
The Cobra Positioners are mounted in an optical bench at the focal plane and are placed
8 mm apart to create a 2,400-positioner field in a hexagonal pattern to allow efficient tiling of thesky region of interest. Each positioner has a circular patrol region of 9.5 mm; this, combinedwith the 8-mm spacing, ensures complete coverage of the populated area of the focal plane. The
populated area fully covers a 1.3-degree field over the area of the focal plane with higher
throughput, resulting in more efficient use of the positioners (Figure 3.10-3).

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1.0

Optical Coupling

0.9

0.8

COUPLING EFFICIENCY

0.7

Complete Coverage with Positioners

0.6

35%
Coveragewith
Positioners

0.5

0.4

0.3

Nominal
Field
Edge

0.2

0.1

0.0
0.000

0.100

0.200

0.300

0.400

0.500

0.600

0.700

0.800

0.900

FOCAL PLANE SEMI-FIELD (Degrees)

Figure 3.10-3: Optical coupling efficiency versus distance from the center of the field of view.
The plot shows the field coverage of the positioner array. Coupling includes the optical
performance of the WFC and the numerical aperture of the fiber, but does not include seeing or
WFC internal flexure.

The actuators are controlled in an open-loop mode with position information obtained after
the fiber field is moved. Their positions are verified by back illuminating the optical fibers andmeasuring its location using a CCD metrology camera.
The key to achieving rapid and accurate target acquisition are the Cobra Fiber Positioners,
Figure 3.10-4. Cobra’s small size, high speed, and high resolution make it a key component of
WFMOS. The critical technology is a new class of rotary motors, called Rotary Squiggle Motors, which were developed specifically for WFMOS by a small team of engineers at JPL and
New Scale Technologies (NST). These motors set new benchmarks for small size and resolution
performance which has not been achieved in electromagnetic motors. This simple, robust rotary
piezo tube motor has been scaled to 3.3 mm diameter, without significant loss of power efficiency. JPL’s experience in high reliability spaceflight mechanisms combined with New Scale
Technologies existing linear motor design and analysis skills as well as mass production capabilities is the perfect teaming relationship for the needs of WFMOS. The teams collaborated in a
rapid 8 month development of the Cobra Positioner that successfully surpassed the requirementsfor positioning accuracy and speed.

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Figure 3.10-4: The Prototype Cobra Fiber Positioner

The motive force of the Rotary Squiggle Motor is provided by two pairs of piezo plates
mounted on a flexible tube that are actuated in a coordinated “hula” fashion to produce rotary
motion (Figure 3.10-5). The Cobra Positioner is made up of a 2.4-mm motor and a 4.6-mm motor, which match the sizes required to ensure the WFMOS survey speed.

Figure 3.10-5: Rotary Squiggle Motor Motive Force Schematic

The Cobra unit is a 7.7-mm diameter, theta-phi style, 2-axis fiber positioner, with the dependant axis offset from the independent axis, enabling the optical fiber to be placed anywhere ina 9.5-mm patrol region, as shown in Figure 3.10-6.

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Figure 3.10-6: Patrol Region of a single fiber positioner (top view). The two motors in eachpositioner have axes that are offset from one another, allowing the fiber tip to travel anywhere
within the patrol region.

Features have been included in the Cobra Positioning system to make integration and test
more efficient. The Cobra Positioners are fabricated and assembled with tolerances that keep thepatrol areas sufficiently parallel to the focal surface to eliminate all but one manual adjustment
during assembly. Also, the positioners are divided up into separable modules to make integration, testing, and servicing easier. Each module has its own electronics module, making verification of mechanical and electrical function and performance possible early in the construction
phase. These features are in line with the overall project strategy to allow early and complete
testing of subsystems.
The Cobra Electronics uses a distributed-driver approach. This approach facilitates maximum
motor movement flexibility with minimum wiring. A block diagram of the approach is shown in
Figure 3.10-7. The Positioner System key component is the Positioner Module. This Module
consists of a unique addressable FPGA and a driver per motor in the module.
The basic approach has three steps:
1. Motor drive information is created in the positioning computer from data generated by
the Metrology System. This consists of the Unique Module ID, the motor number in themodule, the duration of the motors operation, and the direction CW or CCW of the motor.
2. This information is transmitted through an optical Ethernet connection to the Ethernet Interface, which reconverts the Ethernet Data format back into data commands that are
forwarded to the proper Positioner Module FPGA.

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3. The FPGA generates the signals that have the proper duration frequency and phasing tocause the motor to rotate properly.

Figure 3.10-7: Positioner System Block Diagram

Proper packaging design and techniques are necessary to make the 2.5-volt signals from the
FPGA compatible with the 120-volt square wave out of the drivers. To accomplish the electronics packaging, chip on board (COB), automated assembly, and automated testing techniques will
be utilized. Testing will be performed to ensure noise and cross-talk issues are addressed. Sufficient resources have been allocated to allow early testing (by Positioner PDR) and multiple design cycles to characterize the performance of the electronics and bring the design to maturity.
Once assembled and verified, the Cobra modules are then integrated to the Cobra Optic
Bench, which interfaces to a tilt/rotation test fixture via the Positioner Alignment System. Whenthe entire bench is assembled and the electronic modules are daisy chained together, the bench isfunctionally validated. The deformations due to temperature and gravity are acceptance tested
before the components are integrated and verified at higher levels of assembly.
Integration of fibers into the modules can be done in parallel. The design and schedule allows
for integration of the fibers through the center of the positioners; this is a configuration that has
been tested and meets the throughput requirements. However, schedule and workforce can berecovered by routing the fibers externally to the positioner. Early testing will determine whether
the throughput requirements are met. If the requirements are met, the integration approach can be
simplified and accelerated.
The WFMOS control system provides the positioner movement planning software (MPS)
with the target and the current fiber locations—as measured by the WFMOS Metrology System—of the individual optical fibers, specified in the coordinate system of the WFMOS instrument measurement plane. The MPS computes the motor drive information and generates an output file consisting of the sequence of movements for the embedded controller software to40

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command each positioner motor to move the fibers to the target position.
After each sequence of maneuvers, the MPS will receive an updated map of the positioner fiber locations and generate a new command sequence to reduce the error between the target andactual positions. The software includes a learning algorithm to incorporate the history of movements of each positioner with respect to parameters such as starting position, length of move, or
temperature in planning the next move. The Cobra Positioners have a 9.5-mm patrol region, but
are spaced only 8 mm apart. This overlap ensures complete coverage by the positioners, but alsomeans that adjacent positioners may collide if their sequence of maneuvers is not carefully coordinated. When computing the sequence of motor commands for an individual Cobra positioner,
the MPS considers current locations of the six adjacent positioners with which its patrol region
overlaps. The simple solution is to run each positioner radially when transiting common patrol
areas to avoid collisions.
A prototype testbed was constructed to demonstrate the performance of the Cobra Fiber Positioner. The testbed was set up with a back illuminated fiber and a metrology feedback camera
system analogous to the system being proposed for WFMOS (Figure 3.10-8). Even though thiswas a rapid, low-cost development, this first unit was able to position within 5 μm within six iterations, 90% of the time.
Although this performance exceeds the requirements, mechanical and software control improvements have been identified in anticipation of performance degradations not captured in the
prototype test. Differences in large scale production, driver design, subsystem communication,
lifetime, neighboring positioner collisions, and environment all may have effects on the numberof iterations to achieve convergence. We will make improvements in the design where theycome at little or no cost. Optimization of torque versus resolution, the rotor-stator interface design, control algorithm maturation and improvements in manufacturing and assembly techniques
will lead to better resolution, greater reliability, and/or reduced manufacturing costs and gives
confidence the requirements will be met.

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Figure 3.10-8: Prototype Positioner Testbed. A light source illuminates a fiber optic mounted in
the Cobra positioner unit under test. A CCD camera accurately measures the position of thefiber tip.

3.11 Fiber System
The function of the Fiber System is to collect light from the prime focus and deliver it to thethree spectrometers located off the telescope. Based on science needs, the optimum fiber is a
step-index, silica fiber with core diameter of 107 μm, cladding diameter of 140 μm and numerical-aperture (NA) of 0.23. Several commercial vendors, such as Polymicro Technologies, Nufern, and OFS/Furukawa, are willing to provide such custom fibers. Preliminary fiber-routing design indicates that 60-m long fibers would be needed from the prime focus unit to thespectrometer. Figure 3.11-1 shows key elements of the fiber system. Though the fibers, and theapplication they are being used for, may be unique, it behooves us to use technology and processes developed by the telecom and fiber-optics industry. We will, therefore, rely heavily on industry standards and practices to build cables and connectors; we will rely on custom fabrication
only where appropriate.
From a system performance view, the key requirement on the fiber system is its throughput.
Based on preliminary design, the fiber system throughput ranges from 0.4 at the blue end of thespectrum to 0.65 at the red end of the spectrum (see Figure 3.11-2). Three dominant factors affect the throughput of the fiber system: 1) attenuation, 2) focal ratio degradation (FRD), and3) connector losses.

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WIDE FIELD FIBER-FED OPTICAL
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Figure 3.11-1: Schematic of the fiber system from the positioners to the spectrograph. Picturesof key hardware from risk-reduction activities are also shown.

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Figure 3.11-2: Fiber Throughput and the Contribution from Various Components

The bulk attenuation of low-OH silica fibers, typically specified as loss in dB/km, is highly
wavelength dependent in the 400–1,000–nm band. This bulk attenuation is well characterized
and is practically independent of vendor or core size.
To facilitate WFMOS installation and maintenance, fibers need to be sectioned and arranged
for interconnection, which in turn introduces insertion losses. To connect up 2,400 fibers in areasonable time, multi-fiber connectors are necessary. Fabrication of custom very-high-density
fiber connectors with 30 rows of 28 fibers (rectangular pattern with fiber pitch of 200 μm) areplanned to support WFMOS. To support 2,400 fibers, three such mating connector pairs will befabricated, and the extra fibers will be used as spares or for alignment monitoring. LNA has built
and tested a 5 × 5-fiber connector as part of preliminary risk-reduction activities. Test results indicate that the custom connectors meet requirements and that they can be scaled to the desired
fiber counts. In the current baseline, there will be just one connector in each cable, and thethroughput budget assumes an achievable connector efficiency of 88%.
To minimize insertion loss degradation with repeated mate/de-mate cycles, we plan to use adynamic connector scheme, in which one connector (mounted on a 3-DOF stage) is activelyaligned to the other connector during installation. Tests on prototype indicate that positioning
accuracies of <1 μm are achievable. Before settling on the current design we evaluated available
commercial connectors. Readily available MPO/MT connectors support as many as 72 fibers andprovide fiber positioning accuracies of 3 μm with mechanical guide pins. For a 100-μm core
MMF, insertion losses are expected to be below 0.4 dB (or better than 90% throughput) with adegradation of <0.2 dB insertion loss after 500 mate/de-mate cycles. Unfortunately, most of thecommercial products (including MPO/MT connectors) only support 125-μm cladding fibers, andnot the 140-μm fibers baselined for WFMOS. One of the descope options available to us is thus
the use of commercial products along with 105/125-μm fiber.
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Focal ratio degradation (FRD), the third major source of throughput error, is a measure of
how well the fiber preserves the angular distribution of the input light. FRD is often hard toquantify and depends to a large extent on fiber geometry (core diameter and core/cladding concentricity) variation and details of the routing (bends, twists, stresses on fiber). Proper NA
matching between the telescope, fiber, and spectrometer is critical to avoid losses from overfilling or under-filling the fiber or the spectrometer. Fiber mismatch at the connector interface canalso introduce FRD. LNA has devoted substantial effort to understanding and minimizing losses
from FRD. Figure 3.11-2 shows some early results.
Besides the above major contributors there are other loss terms. Fresnel losses (~4%) at bothends of the fiber can be substantially minimized (to ~1%) with broadband anti-reflective(BBAR) coatings. This is relatively easy to do at the slit-block end. AR coating fibers at the positioner end presents logistical challenges; therefore, those ends will be left uncoated. There arealso losses (evanescent field decay) from bending and twisting of fibers. We have allocated 5%
losses due to these effects.
The fibers at the positioner end are placed and bonded in a custom ceramic ferrule that ismounted to the actuator. As the fiber is moved, it twists and coils, either of which can cause
FRD. To understand this, 15 mock-ups of the actuator assembly were built to simulate articulation of fiber in the same manner as the proposed design. Test results indicate that twisting does
not appear to induce an increase in FRD.
To present light to the spectrometer, 800 fiber tips will be lined up in a linear array with apitch of 230 μm. Each “slit block” will contain 80 fibers, and 10 slit blocks will be assembled
end-to-end to build one “slit” for each of the spectrographs. In risk-reduction activities, 80 rowsof fibers were epoxied (with EPOTEK 301) to quartz blocks. Several PTFE blocks were used tohold the position of the fibers accurately during epoxy curing. Once the fibers were bonded, the“slit block” was lapped and polished. For WFMOS, an additional cover plate, with BBAR coating on one side and index-matching gel on the other side, will be placed against the slit block
surface to minimize Fresnel losses.
To protect the fibers from handling and environment, they are enclosed first in a protective
tube and then are grouped into the segmented tubes which act as a protective conduit. The filled
volume of the segmented tubes is calculated to prevent any crushing of the fibers.
As mentioned earlier, the fiber cables are separated into two sections:• The short, approximately 5-m, cables go from the positioners to the top of the PFU
through a strain relief box. The strain relief box is designed to protect the fibers from the
motions of the telescope. This short section stays with the PFU and instrument.
• The long, approximately 55-m, cables run from the connectors on top of the PFU all theway to the spectrograph located in the adjacent room. Part of this long cable—from thespectrograph to the telescope structure—is permanent.
The remaining section is connected to the PFU during instrument operation and stored away(tied to the telescope structure) when not in use.
3.12 Spectrograph
Our WFMOS spectrograph subsystem design comprises a set of three identical generalpurpose spectrographs capable of carrying out all four required surveys (see Section 2) as well as
offering the flexibility for many other types of observation. The spectrograph requirements andbasic parameters are summarized in Tables 3.12-1 and 3.12-2.

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SUMMARY REPORTThe WFMOS surveys demand a large multiplex gain and large wavelength coverage. This,
combined with the 8.2-m Subaru aperture, demands a very large A-omega, the product of thearea of the light beam and its solid angle as measured at the detector. The A-omega of an instrument is an important figure of merit, defining both the information-gathering power and the difficulty of the optical design and, hence, its cost. A-omega for WFMOS is an order of magnitudelarger than that of current large-survey MOS systems (2dF, AAOmega and SDSS) and comparable to that of the multi-IFU MUSE for the VLT.
We have been able to obtain the required A-omega with three spectrographs that have f/1.07
cameras and CCD detectors of 93 mm × 93 mm. We are proposing to use custom CCDs from
LBL. The proposed design (see Figures 3.12-1 and 3.12-2) uses Schmidt optics and meets all
survey requirements. It has a large beam diameter (500 mm) to minimize the central obstruction
losses, give high spectral resolution (as required by the GAHR survey) and keep the field angles
within the camera within sensible limits to maintain good image quality. Schmidt optics is themost cost-effective approach, providing the greatest A-omega per dollar. In our design, Volume
Phase Holographic Gratings (VPHGs) are used in transmission mode to maximize throughput
and versatility. Our design allows spectral resolutions of up to R ~30000.

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Table 3.12-1: Spectrograph requirements based on the four surveys. The first three rows are arepeat of the science case. The other rows are derived requirements.
Requirement

DELZ

DEHZ

GALR

GAHR

Blue

Red

Min λ (nm)

630

420

480

818

480

Max λ (nm)

970

670

544

885

680

Resolving power

3500

1500

5000

5000

20000

SRE width (Angstroms) *

2.82

3.63

1.03

1.70

0.29

Number of SREs per spectrum

1488

688

680

412

6897

Number of spectral orders per fiber

1

1

1

1

4

Length of a single spectral order (mm)

80

37

37

22

93

No. of spectrographs required

2.6

1.2

1.2

0.7

3.0

% of detector area used

86.3%

39.9%

36.3%

22.8%

100%

* SRE – Spectral Resolution Element

Table 3.12-2: Basic derived spectrograph parameters. These are driven by the GAHR survey,
which is the most demanding of the four surveys.
Detector format for one spectrograph

93 mm × 93 mm

Format of a single CCD

93 mm × 46.5 mm, buttable along the long side

Number of spectrographs

3

Collimator f-ratio

f/2.12

Camera f-ratio

f/1.07

Minimum beam size

~300-mm diameter

Fiber pitch at slit

230 µm, center to center

Projected fiber size at detector

54-µm diameter

SRE* width (µm at detector)

54.0 µm

SRE height (µm at detector) = projected fiber pitch

116.25 µm* SRE – Spectral Resolution Element

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Figure 3.12-1: The optical design of a single WFMOS spectrograph, shown here as configured
for the GAHR survey. The diameters of the major optical elements are shown in millimeters.

Figure 3.12-2: CAD model of one of the three identical spectrographs, shown here asconfigured for the GAHR survey. The table is approximately 5 m × 3 m.

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Changing survey modes requires the following:• Changing the camera–collimator angle. The spectral resolution R is directly related tothis angle, with the lowest R corresponding to the straight-through position.
• Changing the VPHG by moving the active VPHG out of the beam on to the selectorslide, moving the slide to the position for the new VPHG and then moving the new
VPHG in to the beam (see Figure 3.12-3).
• Changing the band-pass filter by rotating the filter wheel.
This process is fully automatic and can be done in less than one minute.
The WFMOS design uses VPHGs for all four surveys. VPHGs have several advantages over
the ruled gratings traditionally used in astronomical spectrographs. VPHG peak efficiency is inexcess of 90%. For both the DELZ and DEHZ surveys, the gratings will be used in first order.
For the two Galactic Archaeology surveys, higher orders are used.
The design is extremely versatile and general-purpose. The survey gratings can potentially be
used at other wavelengths and resolutions by changing the camera–collimator angle. For example, the GAHR grating gives R ~30,000 at the maximum camera–collimator angle, and any
wavelength sub-range (about 30-nm long) in the 420-nm–970-nm range can be selected by selecting the appropriate order with a band-pass filter. The band-pass filters are small and inexpensive. A great advantage of a VPHG spectrograph with a variable camera-collimator angle is thatthe gratings can always be used so that they satisfy the Bragg condition and, therefore, work attheir peak efficiency. The design has additional flexibility in that it is possible to install new gratings for future observations.

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Figure 3.12-3: A time sequence showing the system reconfiguring between the DELZ mode
(top left) and the GAHR mode (bottom right). Note that camera, the filter wheel and the gratingexchanger all move at the same time.

3.13 Detector System
A mosaic of two close-butted CCDs will be located at the prime focus of the reflective f/1.07
spectrograph camera of each of the three WFMOS spectrographs. Four hundred spectra will be
projected on each CCD with fibers spaced at the slit to avoid the small gap between CCDs. A
plano convex field flattener in close proximity to the CCDs will serve as the window to a very
compact vacuum housing surrounding the CCD mosaic connected by a slender duct to a low vibration Joule-Thompson cooler placed outside the spectrograph beam. The design minimizes
beam obstruction on axis and is tapered to avoid additional beam obstruction off axis. Since thespectrograph is gravity-stable and vibration isolated, the CCD cabin is supported only by the rectangular duct. The entire vacuum enclosure with the exception of a Kovar window frame will be
constructed from stainless steel, electropolished and cold coated on the interior for low emissivity. The exterior finish will be matt black and have a stepped profile to eliminate most grazing
incidence reflections. Image ghosting is minimized by the very low reflectivity of the multi-layer
coated CCD, by AR coatings on the field flattener, and by spacing the field flattener 1.7 mmfrom the CCD. The use of stiff materials and minimization of cantilevered mass pushes the lowest resonant frequencies to 27 Hz and 83 Hz and 184 Hz. These figures can be significantly improved by further tuning of the design so that significant attenuation will be achieved by thespectrograph’s vibration damping mounts, which will have 2 pole roll off above 1.5 Hz.

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Figure 3.13-1: The CCDs will be located within an evacuated enclosure supported by a narrowduct carrying a cold finger from a Joule Thomson (low vibration) closed-cycle cooler outside thespectrograph beam.

CCD columns will be oriented parallel to the dispersion direction to allow faster region-ofinterest readout for DEHZ and GALR surveys, which use only half the detector area. This orientation also supports rapid binning of spectra to just a few lines to support rapid measurement of
total flux during raster scans for calibrating positioner coordinates.
A custom CCD will be procured, since this offers significant CCD performance and opticaladvantages. The optimal CCD size for good optical performance and minimum spectrograph costis about half way between the standard large-format CCD size (61 mm) and a mosaic of two
such CCDs. To minimize the speed of the camera and thus improve optical performance, theCCD size has been increased to 46.5 mm by 93 mm, the maximum allowing two CCDs to fit into
the available 138 mm diameter on a 150 mm silicon wafer. The 16 μm pixel size has been optimized for the best compromise between resolution (for high signal-to-noise GA surveys) and
S/N for the DE surveys where exposures are shorter and read noise has some impact. This leads
to CCD dimensions, 2,880 columns and 5,760 rows, with 3 to 3.5 pixels across the FWHM of theimage of the fiber after aberrations and lateral charge diffusion are taken into account.
Quantum efficiency is projected to be very high over the entire 420-nm–970-nm wavelength
range (Figure 3.13-2) and compares favorably to competing CCDs, particularly at the extremes.
This remarkably high QE will be achieved by combining several recent, but proven, innovations:
1. Thick (450 μm), fully depleted p-channel CCDs, designed by LBL and fabricated by
DALSA, will substantially enhance the QE at the longest wavelengths. Full depletion forsuch thick CCDs has been made possible by LBNL’s now mature design, which supports
particularly high backside bias voltages (tested to 220 V). Notwithstanding the greater
thickness, lateral diffusion and defocus with absorption depth can be shown to be negligible compared to the size of the fiber image.

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2. The delta doping process developed by JPL specifically for these p-channel CCDs will beapplied to eliminate the absorption in the backside contact, providing reflection-limited
blue response not otherwise seen shortward of 500 nm.
3. JPL will deposit a custom dual-layer sputtered TiO2/SiO2 coating to reduce reflectivity toless than 4%, from 420 nm to 970 nm.
4. The low dark current seen in the delta doped LBL detectors allows them to be operated at–100°C or warmer, providing enhanced red QE due to the shorter absorption length athigher temperatures.

Figure 3.13-2: Theoretical 1-R and lower bound on QE compared to next best competitor. The
Hamamatsu can be improved in the red by increasing thickness but only to 300μm, and by
applying the Ti02/SiO2 coating designed by JPL.

LBL will perform cold wafer probing prior to backside processing and packaging. Packaging
by JPL will assure less than 10 μm height variation across the mosaic of two detectors, whileproviding for small gaps between each CCD and its neighbor (or vacuum housing). Caltech will
characterize and integrate the CCDs into the cameras.
Full-frame readout time will be reduced to 16 s to match the fast reconfiguration time of the
Cobra fiber positioner, while also reducing overheads between multiple exposures on the same
field to <5%. This will be achieved by reading through 20 amplifiers per CCD to compensate forthe low pixel rates (~50 kHz) required to achieve very low noise (Figure 3.12-3).
Yield loss due to the high amplifier count will be concealed by providing 40 amplifiers per
CCD and selecting 20. An amplifier will be located at each end of each of 20 serial register segments. The best amplifier will be selected for each segment by setting by clock and video jumpers to determine the serial clocking direction. In addition to improving yield, this will tighten thenoise distribution, eliminating outliers, and will reduce the mean noise.

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Figure 3.13-3: Read noise versus dwell time for three DECam CCDs manufactured by
LBNL/DALSA. Courtesy of Juan Estrada, Fermi Lab.

Fabrication and testing of 84 CCDs has been budgeted to produce 6 science-grade devices,safely more than the yield experienced by the 62 CCD Dark Energy Survey Camera. Images will
be dithered in the spatial direction so that cosmetic defects will be removed along with cosmic
rays, allowing specifications for cosmetic defects to be relaxed. Three hundred-μm thickHamamatsu CCDs without delta doping (but possibly with JPL’s multilayer coating design) are
offered as a technology fallback option, although this route appears to be more expensive (unlesswavelength coverage or multiplex advantage is sacrificed) and delivers lower performance atboth wavelength extremes.
3.14 System Software
The WFMOS software is responsible for:• Preparing and managing WFMOS observations in advance.
• Selecting the observations to be made each night.
• Top-level control of WFMOS observations at the Subaru observatory.
• Providing engineering and test software.
• Batch processing of WFMOS data in real time to monitor survey data quality.
3.14.1 Observation Preparation
WFMOS will be capable of observing 2,400 science targets simultaneously in a single observation and, in total, a given WFMOS survey will yield data for up to 4 million science targets.
Observers need a means of automating the generation of their surveys. The WFMOS observation
preparation software will provide the observer with facilities for selecting the science targets thatmake up a WFMOS survey and arranging those targets into a collection of schedulable observations. It will also provide configuration software for matching targets efficiently to fibers. The
end product of the observation preparation software is a database of observation commands waiting to be scheduled.
3.14.2 Observation Management
The WFMOS software provides facilities for managing prepared observations. Priorities canbe assigned, extra science targets can be added to observations that have spare fibers, and obser-

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vations can be modified to bring them up to date with any changes in instrument capabilities.
Nights of WFMOS observations are scheduled by the Subaru observatory.
3.14.3 Observation Execution
The WFMOS control software runs at the Subaru observatory and interacts with the local
computer infrastructure through Subaru’s SOSS software. Descriptions of the WFMOS observation descriptions required for a particular night are extracted from the database created by theobservation preparation software and may be executed by the observer or observatory staff
member.
The WFMOS control software is responsible for coordinating all the underlying WFMOS
subsystems. The control hierarchy and block diagram are shown in Figures 3.14-1 and 3.14-2,
respectively. The control software manages the sequence of operations needed to make each observation, acquire the data and save it to the data store. A typical sequence would be:• Acquire the target field with the telescope and configure the spectrograph.
• Use the feedback provided by the metrology camera to configure the fiber positioner.
• Acquire guide stars using the autoguiders.
• Acquire data with the science detectors.
• Collect and store FITS header information.
• Send the data to the data pipeline.
The sequencing steps are overlapped wherever possible to maximize the survey efficiency;
for example, the next target can be acquired while the science detectors are still reading out.
The WFMOS control software will manage the subsystem interfaces shown in the block diagram and ensure all the interfaces meet the required standards of maintainability. A Pythonbased communication infrastructure will be provided.
3.14.4 Engineering and Test Software
Engineering software will be used to set up and calibrate the instrument during integrationand testing, and when WFMOS is installed on the Subaru telescope. The end product of the instrument calibration software is an instrument description that is shared with all the WFMOS
software systems. The engineering software will allow the instrument description to be kept up
to date; for example, if any fibers or positioners break or if the spectrograph has new components
installed.
The WFMOS software will provide diagnostic and testing facilities that can be used throughout the development and testing phases of the instrument. Automatic test scripts will ensure that
instrument components are tested frequently and rigorously. The engineering facilities will remain after delivery to provide diagnostic tools for investigating any faults and problems.

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Figure 3.14-1: Control Hierarchy
Science

Guide Stars

Light:
Science
Guide
Calibration

STARS Archive

Subaru Telescope

Calibration
lamps

Controlsignals

SOSS

Subaru TCS

Data

ObserverPHSC

WFMOS Control System

PHSC hexapod
Instrument rotatorWFMOS Data
Handling and
Pipeline System

WFC + ADC

Autoguider
Metrology Detector
Positioner

Fiducial
lamps

Slit block 1
Positioner lamp 1

Slit block 2
Positioner lamp 2

FibresSlit block 3

Positioner lamp 3Spectrograph 1

Science Detector 1

Spectrograph 2

Science Detector 2

Spectrograph 3

Science Detector 3

Figure 3.14-2: Block Diagram. Subaru infrastructure is shown in green, Subaru-suppliedcomponents are shown in cyan, and WFMOS-supplied components are shown in yellow.

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3.15 Data Reduction
The Data Reduction (DR) Subsystem comprises four basic sub-components:• A data handling system (DHS).
• An infrastructure for running DR.
• Quick-look data visualization.
• Data reduction routines.
The subsystem directly interfaces to the Detector Controller(s) (to whom it provides data
storage facilities) and to the Software Control subsystem, which provides overall control andhandles all external interfaces.
The bulk data interface will be provided by Network Attached Storage (NAS) system; thedata store will be directly available through mount points to the other systems. Control interfaces
will be provided as Python wrappers; in this way, the other systems can be insulated from theunderlying communication system. The underlying communications will be handled either by an
ORB (if the design re-uses the Subaru SIMCAM software communications) or by a common
message bus (if the Gemini Pipeline Reduction System - PRS communications are preferred).
Communications with the Quick-look tools will be possible either by direct access to the NAS or
via IVOA Spectral and Image access protocols, thereby allowing any VO-compliant tool
(VOSpec, SPLAT, etc.) to be able to access the data.

Figure 3.15-1: DR Subsystem Block Diagram and Interfacing
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SUMMARY REPORTThe Data Reduction routines themselves have been split into two packages; Package A (standard spectroscopic data reduction that provides “science ready” data) and Package B (that provides scientific quantities, such as red-shift, from the spectra). Package A follows the data flow
shown in Figure 3.15-2:• First the data are amalgamated and identified, and the relevant calibrations selected.
• The image data are flat-fielded and bias corrected.
• The orders are traced on the image, then extracted to form spectral data.
• The spectral data are wavelength calibrated.
• The data are then made available for further DR (Package B), QC checks and archiving.

Figure 3.15-2: DR Data Flow

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4.0

Instrument Performance

4.1

Optical Throughput
Throughput is a key parameter as it ultimately governs how quickly observations can be
made. The Optical Throughput allocation, where each element of the design was assigned athroughput to match or exceed was created and these values are included as requirements in the
FPRD. As requirements, these values represent the minimum performance; the instrument design
strategy is to maximize performance and the Current Best Estimate (CBE) is presented here. As
the instrument design matures, continued improvement is sought through subsystem design
choices and balancing resources throughout the instrument system to maximize total throughput.
The instrument CBE is presented in Table 4.1-1 and is supplemented with the estimatedthroughput from external factors such as the telescope primary, the WFC, and atmospheric dispersions. The instrument performance estimates, also include the variations between the different
surveys. For example, the Dark Energy surveys require a different order sorting filter type fromthe Galactic Archaeology surveys and produce a different throughput estimate for each survey
type. The throughput is presented at the wavelength extremes of 420 nm and 970 nm and several
wavelengths in between.

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Table 4.1-1: Throughput budget for Dark Energy and Galactic Archaeology surveys. The spectrograph efficiency differs for the twosurvey types because of differences in the order sorting filter. The bottom line throughput number for the Galactic Archaeology
surveys reflects this difference. The throughput does not include optical coupling at the fiber entrance as this depends on seeing andthe target (e.g. galaxy size). These factors are accounted for separately in exposure time calculations. The entry for the WFC
vignetting is an average across the field.
Wavelength (nm)
420

440

550

Component

640

790

970

Throughput

Atmosphere

0.83

0.83

0.89

0.9

0.9

0.9

Secondary Obscuration

0.95

0.95

0.95

0.95

0.95

0.95

Primary Reflectance

0.9

0.9

0.9

0.9

0.85

0.85

WFC Throughput

0.88

0.88

0.88

0.88

0.88

0.82

WFC Vignetting

0.83

0.83

0.83

0.83

0.83

0.83

Uncorrected Atmospheric Dispersion

0.995

0.995

0.995

0.995

0.995

0.995

Corrector Plate Throughput

0.98

0.98

0.98

0.98

0.98

0.98

Misalignment between WFMOS/HPFU

0.98

0.98

0.98

0.98

0.98

0.98

Obscuration of pupil by Metrology System

0.995

0.995

0.995

0.995

0.995

0.995

Obscuration of pupil by Spectral Cal System

0.995

0.995

0.995

0.995

0.995

0.995

Fiber Positioners

0.89

0.89

0.89

0.89

0.89

0.89

Fiber System

0.55

0.55

0.61

0.66

0.68

0.68

Spectrograph

0.689

0.689

0.689

0.689

0.689

0.689

CCD QE

0.895

0.931

0.957

0.973

0.997

0.931

DE Instrument throughput

0.29

0.29

0.34

0.38

0.40

0.37

GA Instrument throughput

0.23

0.25

0.29

0.32

0.34

0.26

DE Total throughput

0.15

0.15

0.19

0.21

0.21

0.18

GA Total throughput

0.12

0.13

0.16

0.18

0.18

0.13

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4.2

Reconfiguration TimeIn addition to instrument throughput, the other driver of observing efficiency is observational
overhead. In order to minimize overhead, parallel activities are planned. Immediately after theprevious exposure is complete, we begin reading out the three science detectors. Simultaneously
we begin slewing the telescope to the new target location and repositioning the fibers. This involves illuminating the fibers, imaging the fiber tips with the Metrology camera, computing thecentroids of the imaged fiber tips, and moving the fiber tips. The repositioning sequence is repeated until the fibers are within tolerance of the object coordinates. The time budget allows six
iterations for fiber positioning.
The fibers are back-lit from the spectrographs and are arranged in the focal plane so that notwo adjacent fibers are illuminated from a single spectrograph. This allows unambiguous identification since there are overlapping patrol regions. Table 4.2-1 illustrates the activities happening
in parallel. Once the metrology cycle is complete, the Acquisition and Guide System acquires theguide targets. When the A&G system notifies the Control System that it has a lock, the instrument starts a new exposure. The entire process will take less than 40 s.
Table 4.2-1: Positioner Configuration Timeline
Start Time
0
1
3
5
7
9
11
13
15
17
19
21
23
25
27
29
31
33
35
37
38

Duration (s)1Move A-1
2
Image A-1
2
Compute A-1
2
Move A-2
2
Image A-2
2
Compute A-2
2
Move A-3
2
Image A-3
2
Compute A-3
2
Move A-4
2
Image A-4
2
Compute A-4
2
Move A-5
2
Image A-5
2
Compute A-5
2
Move A-6
2
Image A-6
2
Compute A-6
2
1

Steps being Taken in Parallel
Move B-1
Image B-1
Compute B-1
Move B-2
Image B-2
Compute B-2
Move B-3
Image B-3
Compute B-3
Move B-4
Image B-4
Compute B-4
Move B-5
Image B-5
Compute B-5
Move B-6
Image B-6
Compute B-6

4.3

Move C-1
Image C-1
Compute C-1
Move C-2
Image C-2
Compute C-2
Move C-3
Image C-3
Compute C-3
Move C-4
Image C-4
Compute C-4
Move C-5
Image C-5
Compute C-5
Move C-6
Image C-6
Compute C-6

Read out of Science
Detectors
Move Telescope to
New Position
Acquire Target w/
A&G system

Allocation Efficiency
An important aspect of any positioner scheme is the efficiency with which the individual
positioner elements can be allocated to the targets of interest for spectroscopy. This allocation
efficiency is a measure of how effectively the spectroscopic capability of the instrument is being
used. WFMOS has been designed to provide high allocation efficiency for both the low and highspectrograph modes. If the number of targets is very large compared to the number of positioner
elements, then every element will have many targets available to choose from and the positioner

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allocation success rate will be high. If there are few targets, then most positioner elements will
not have a target available. At intermediate target density, the success rate will be moderate, andwill depend on the degree of overlap of adjacent patrol regions. A full discussion of the allocation efficiency issues and trade-offs appears in the OCDD, Section 6.
4.4

Exposure Times and Survey Time Estimate
As discussed above, WFMOS provides high throughput and high positioner allocation efficiency, allowing the primary surveys to be accomplished in modest telescope time. Table 4.4-1
summarizes the required exposure times and survey times for the primary surveys. More extensive detail of the survey efficiency and exposure time calculator appear in the OCDD.
Table 4.4-1: Exposure and survey times for the low-redshift Dark Energy survey (DELZ), thehigh-redshift Dark Energy survey (DEHZ), the low-resolution Galactic Archaeology survey(GALR), and the high-resolution Galactic Archaeology survey (GAHR).
SurveyResolution

Mag

S/N

Obs (min)

Survey time (hrs)

DELZ

3500

RAB < 23.0

Line = 7

15

832

DEHZ

1500

RAB < 25

Cont = 4

240

286

GALR

5000

V = 20

10–15 / Å

30

765

GAHR

20,000V = 17150 / Å

60

1125

4.5

System Versatility and Upgrades
The conceptual design of the current instrument has been tuned to delivery optimal performance for the Gemini surveys (DE/GA) while attempting to keep costs low. However, as introduced in Section 2, as the unique capabilities of WFMOS will have major impact in many areasof auxiliary science, we have borne this in mind in our design study.
There are significant features in the major elements of this design, the fiber positioner assembly and the spectrograph, that ensure versatility and permit possible upgrades. The key tokeeping upgrade options open is clean interfaces between subsystems and design modularity.
The highlights of potential upgrades are:
• Increase in the total number of positioners—possible because of the design of theindividual positioner and the spectrograph design.
• Increase in spectral resolution of the spectrographs.
• Survey changes with different gratings.
4.5.1 Positioner
The current design has 2,400 positioners arranged in a hexagonal pattern centered on the focal plane that covers 1.2 to 1.4 degree field. This configuration was selected to optimize the datacollected with respect to cost. The outer edges of the 1.5 degree focal plane are affected by vignetting reducing the throughput and subsequently making those fibers beyond 1.4 degrees lessefficient (Figure 3.10-3). However, even taking this vignetting effect into account, a decrease inobservation time can be obtained by fully populating the focal plane with positioners. The number of positioners can be increased to 3,200 with no change in the design of the existing Cobra
positioner, its control system, or its electronics architecture.
4.5.2 Spectrograph
The spectrograph design easily allows an increase in the number of positioners as well asother design upgrades. Since each spectrograph is identical, adding a fourth spectrograph would
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accommodate more positioners. No recurring engineering costs are added.
The spectrograph beam size is 500 mm in diameter, and the range of camera-collimator angles is 0 to 110 degrees. A grating can be designed for any angle in this range, giving a range ofspectral resolutions from 0 to 30,000. The width of the detector in SREs (spectral resolution elements) is constant—there are always 1724 of them. So as the resolution increases the wavelength
coverage decreases. The whole wavelength range 420 nm–970 nm is obtained at R ~2178 at
12 deg, whereas 31.6 nm of spectral coverage in a single order is obtained at R ~30,000 at
100 degree. The spectrographs can, therefore, be used to do other science programs if new gratings are purchased. Furthermore, the gratings purchased for the four surveys can be used at different angles to give different wavelength ranges.
Our system will have significant sensitivity below 420 nm and above 970 nm, which willdefinitely be of great interest for auxiliary science.

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5.0

Integration and Test

5.1

Integration Approach
Our history in delivering science instruments has enabled us to develop a systematic approach to I&T where it is common to have subsystems developed and tested by different groups.
Emphasis on early validation at the subsystem level with clear definitions of interfaces allows
verification of functionality and identification of problems early and minimizes the problems encountered when the instrument is finally assembled and tested as a system. Initial sub-systemlevel validation also provides early test data to confirm interfaces, especially data content andform. The verification process of the instrument begins with the development of the Acceptance
Test Plan (ATP) to define the series of tests, inspection, and analyses used to validate the instrument requirements. This plan will emphasize subsystem level tests and subassembly tests to verify functionality and performance early and to identify any problems and minimize telescope
time for acceptance tests. At each stage we weigh the differences in response or performance
when simulating an interface against the complexity of higher level tests, time and personnel
necessary to debug higher level assembly tests, and finally telescope time at the full system level
to determine when each requirement should be validated.
The strategy in developing the ATP is:• Develop tests with clear objectives and success criteria.
• Validate as many requirements as possible at the lowest level.
• Select representative tests to confirm tests at the previous stage after the next integration
step.
• Select simple interfaces and identify risks associated with use of simulators and support
equipment compared to a higher system-level test.
• Minimize telescope time for Acceptance Tests by identifying the minimal set of tests that,in conjunction with previous tests and analysis, will fully validate the instrument.
An example of how this process works for the PFI is described in Section 2.14.6 of the Detailed Design Document. A similar approach is developed for every assembly and subsystem.
These tests are defined as the design progresses to ensure features needed for I&T are built into
the software and hardware and the necessary laboratory facilities are available at the appropriate
time.
Examination of the instrument block diagram (Figure 3.2-1) shows this instrument can be integrated along two paths: one involving the telescope-mounted hardware and the other for the
spectrograph, detectors and data pipeline. The software control system will be installed andtested at both integration sites. All of this occurs prior to integration onto the telescope, and theonly test missed is the overall optical throughput, which can easily be calculated from individual
tests. These parallel paths will also be followed in the initial telescope installation work before
the final on-sky work commences. The advantages of this approach are minimizing the risk toschedule by allowing parallel independent progress and flexibility for as long as possible andmaximizing the total resource available to the integration Phase by using two integration sites.
The full plan is detailed below and shown schematically in Figure 5.1-1. The System Software will require careful management as modifications and bugs will be appearing simultaneously from two sites/time zones.

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Figure 5.1-1: The Instrument Integration and Test flow follows two parallel paths before final integration onto the telescope:
spectrograph/detector/data pipeline integration and PFI/metrology integration.

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5.1.1 Spectrograph/Detector/Data Pipeline/WFMOS Control Subsystems’ Path
This integration work will be located at the Institute of Astronomy at the University of Cambridge. The main tasks are:• Mechanical Assembly and alignment of the Detector system to the spectrograph optics.
Confirm cryogenic performance of Detector System in Spectrograph.
• Connect Detector to Data Pipeline Software and confirm files/data are being passed.
• Confirm three systems work properly together using the System Software—it is assumed
individual systems will have already tested their interfaces to the System Software.
• Measurement of the optical performance in ‘imaging mode.’• Measurement of the optical performance in all ‘spectrometer modes’ and calibration ofgrating/spectrograph angles.
• Confirm the performance of on-line data reduction against known spectra.
After formal customer acceptance, the hardware will be disassembled where necessary,
packed and shipped directly to the Subaru telescope.
5.1.2 Prime Focus Instrument (PFI) Path
The second path for the instrument integration and test will be located at the Jet Propulsion
Labs in Pasadena. Prior to this phase, the ‘Cable B’ fibers (the shorter section of the fiber cable)
have been integrated into the positioners and the positioning accuracy, repeatability, and timing
have been validated. The primary remaining task for the PFI/Positioner/Metrology loop is to integrate the metrology system using the system software.
The main tasks of integration and testing at JPL are:• Test of the System Software control loop with the Positioner Assembly and Metrology
Cameras. They will be configured to mimic the instrument setup on the Subaru telescopewith the four metrology cameras placed in approximately the same location relative to the
PFI, as they would be on the Telescope (Figure 5.1-2). Two tilted spherical fold mirrors
can then be used to represent the primary mirror of the telescope for testing and tocompensate for the absence of the Wide Field Corrector. The metrology system iscalibrated to quantify error in the position measurements of the fiber tips.
• Final test will be an optical throughput measurement from positioner to Slit Block (atSpectrograph) and will confirm the repeatability of that performance over several
Connector matings.
Once again, after formal customer acceptance, the hardware will be disassembled where necessary, packed and shipped.

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Figure 5.1-2: Test Configuration of the Metrology, Positioner and PFI Loop at JPL

5.1.3 Integration and Verification at Subaru Telescope
Apart from final verification, a key goal of the work at the Subaru Telescope will be tominimize downtime for the telescope. The final test plan will be revised to benefit from lessons
learned during integration of other instruments at Subaru that occurs before WFMOS delivery(e.g., HyperSuprimeCam and FMOS).
Initially there are three parallel activities which are not expected to need telescope shutdown:
1) re- assembly and test of the spectrographs in their assigned room in the dome, 2) integration of
the System Software and Data Reduction Pipeline, and 3) re- assembly and test of the PFI elements followed by assembly to the PHSC to make a Prime Focus Unit. System Software Integration can probably be done both at sea level or at the summit of Mauna Kea. We expect all spectrograph-related hardware to be delivered directly to the Telescope site. The integration site forassembly of PFI elements into the PHSC will require negotiation with Subaru. These three activities are shown in the flow diagram and include visual inspection, functional checks (eachsubsystem will have the capability of performing a self diagnostic test) and a limited repeat of
the acceptance test which closed out work prior to shipment.
Two other activities occur during this period, namely preparing the telescope structure forand fitting of fiber Cable A (the longer section of the fiber cable) and the Metrology Cameras.
These ought to be possible with daytime engineering only.
It is at this point that the two integration paths finally join together as the Prime Focus Unit isinstalled on the telescope and the fibers are connected to the Spectrograph. This is when coordination with the telescope staff is crucial to develop procedures and schedules which optimize thetest flow but minimize telescope down-time and cost. The key tasks of this final phase are to:• Validate operation of metrology camera using the primary mirror, WFC, and PHSC toposition the fibers.
• Validate operation of A&G cameras on sky.
• Demonstrate and validate target acquisition on sample fields.
• Verify on-sky tracking performance.
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SUMMARY REPORT• Verify optical throughput and measure sensitivity.
• Verify data pipeline on real targets.
• Verify software integration to the telescope.
• Final customer acceptance test.
The instrument is then ready for science commissioning and the main contract would becomplete.

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6.0

Operations

We have designed the WFMOS instrument and science surveys for efficient, streamlined operation. There are several components of operations that are necessary for realizing the scientificreturn of WFMOS. The “Package A” data reduction pipeline is included in this construction proposal and its associated costs. This pipeline produces calibrated spectra, and is described in Section 3. In this section, we summarize the additional components of operations required to complete the Dark Energy and Galactic Archaeology surveys. Additional detail on the scientificsurveys and operational modes of the instrument appear in the OCDD; additional technical detailon the instrument appears in the Detailed Technical Design. The components of operations fall
into the following categories:• Construction of data reduction “Package B”• Conducting the DE and GA science surveys• Installation and removal of WFMOS from the Subaru Prime Focus.
We examine each of these categories in turn.
6.1

Data Reduction Package B
The Package B pipeline use the calibrated spectra from Package A and produces automated,high level science products ready for further analysis. For the DE survey, this consists primarilyof extracting redshifts from the galaxy survey through the identification of the [OII] feature forthe low redshift survey, and through continuum fitting for the high-redshift survey. Galaxy typing will be done for the better s:n targets. We will use a small team of scientists together withprofessional programmers to produce the DE portion of Package B. Associated costs appear in
the Management and Cost volume. The PI delegates scientific management of the effort to Professor John Peacock.
For the GA survey, Package B includes extraction of velocities and metallicities (low resolution) and abundance determination (high resolution). Tasks include 1) automated estimation of
global stellar atmosphere parameters, 2) automated continuum estimation, multi-component line
fitting, and abundance analysis from line lists, 3) automated radial velocity estimation, and 4)
using optical plus near-IR photometry (and distances if available, e.g., from Gaia) to produce
model atmosphere fits to spectra. Associated costs appear in the Management and Cost volume.
The PI delegates scientific management of the effort to Prof. Mike Irwin.
Some of the Package B tasks have significant overlap between the DE and GA surveys.
These include:• Optimizing sky subtraction methods.
• Scattered light modeling using blocked-off fibers.
The apportionment of these tasks will be made at the beginning of the Package B effort.
6.2

Conducting the Dark Energy and Galactic Archaeology Science Surveys
The two primary science surveys are described in detail in the OCDD. Each survey comprises the following steps:• Preparing observations.
• Monitoring survey progress.
• Analysis of results and publications.
Preparing observations involves selecting targets from existing catalogs, optimizing fiber al68

WIDE FIELD FIBER-FED OPTICAL
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SUMMARY REPORT

location efficiency, evaluating survey completeness, deciding tiling strategy, and prioritizing observations considering sky brightness and seeing. An observing tool that allocates targets to fibers will be provided as part of the system software during the construction phase. The additional
steps described here involve interacting with this tool and performing any necessary additional
optimization.
The survey preparations and survey monitoring have many common elements in the two
primary surveys. A number of scientists will participate in the survey phase of WFMOS, and theprimary surveys will be conducted over the course of approximately five years (depending on the
yet-to-be-negotiated agreement with Subaru). The bulk of the day-to-day effort will be conducted by post-docs under the direction of members of the current science team. The estimatedcost for this phase appears in the Management and Cost volume.
6.3

Installation and Removal
We anticipate that WFMOS observations will be interleaved with observations using other
instruments, including the HyperSuprimeCam. Removing WFMOS from the telescope will necessitate the removal of the Prime Focus Unit from the telescope, disconnection of the fiber connector, storage of a portion of the fiber optic cable on the side of the telescope, and possible removal of the metrology cameras. Removal of WFMOS components from the Prime Focus Unit
will be required to allow the Prime Focus Unit to be reconfigured with HyperSuprimeCam. Installation of WFMOS is essentially the reverse of this process. We estimate the Prime Focus Unit
can be removed from the Subaru prime focus in less than three hours and be accomplished entirely with Subaru personnel in a manner similar to standard top-end unit exchanges.
The initial installation of WFMOS components in to the Prime Focus Unit is described in thePrime Focus Instrument section of the Detailed Technical Design. Routine swaps of HSC andWFMOS will be a streamlined version of this process. WFMOS is designed to allow the swap of
HSC and WFMOS components within the PFU to take approximately one full work day. Thisestimate depends sensitively, however, on the detailed HSC design. Two members of the
WFMOS instrument team will be required to supervise the first two swaps. It is also prudent tobudget an additional few days for the first two swaps. Subsequent swaps would be accomplishedentirely with Subaru Observatory personnel. The detailed procedure and timeline for this process
will be developed before the WFMOS preliminary design review.
6.4

Additional Considerations
In addition to the Dark Energy and Galactic Archaeology survey observations proposed here,
there will be classically allocated non-survey observations. Although not formally requested to
do so by Gemini, we believe it prudent to consider the operational scenarios for these potential
additional observations. We envision that many of these surveys will have long-term status andbe queue scheduled. The rest will be a broad range of observing programs proposed by PIs fromwithin the Gemini and Subaru communities. These could also be queue scheduled, or the observers may be given fixed nights and visit the telescope. WFMOS can accommodate these observations with no additional software cost by using the System Software observing preparation tools
and the Package A data reduction pipeline. The operation of WFMOS will not be very different
from that of other fully automated multi-fiber spectroscopic survey instruments, such as FMOS,
2dF+AAOmega, and FLAMES.

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7.0

Management and Cost

7.1

Management and System Engineering
Instruments on the scale of WFMOS have few precedents in the domain of ground-based
astronomy. Indeed, the cost and complexity of WFMOS are more typical of space-based
astronomy instruments and missions. As such, the management approach to WFMOS should beappropriate to the complexity of this development and draw on lessons learned from the
management of other developments of this scale
The WFMOS team has chosen the Jet Propulsion Laboratory to be the lead partner institution
to manage and coordinate these activities. In this role, JPL provides overall project management
and project systems engineering in addition to its technical contributions to WFMOS. JPL has
broad experience in project management on a range of projects for both NASA and non-NASAprojects and has delivered over 40 spaceflight, airborne, balloon, and ground instruments in the$10 M–$100 M range over the past 20 years. Based on this experience, JPL has identified practices that are applied through the project phases to ensure that WFMOS will successfully deliver
within schedule and budget:
1. Detailed planning during pre-project phases, including detailed schedules, budgets, andresource requirements.
2. Identification of cost containment strategies, including descope lists.
3. Identification of potential risks and mitigation strategies.
4. Provision for adequate budget and schedule reserves.
5. Clear statements of roles and responsibilities for key project members (PI, PM, PSE, PS)
and decision-making.
6. Early identification of project communication and coordination strategies.
Managing an instrument developed by multiple international partners requires additional effort and is routinely done at JPL. This requires detailed knowledge of the development of the instrument in order to correctly assess progress and identify trends in schedule or cost growthquickly to take corrective action. Experience has shown the key to managing this type of projectis through continuous communication with all partners, focused attention on the cost and schedule, and adherence to requirements and ICDs for design structure. JPL is committed to delivering
this instrument on time and at cost and, during the development, providing complete insight on
the instrument progress to Gemini. This is done through the monthly reports that include technical status, milestone status, cost and schedule detail, and any issues and mitigation.
A recent example of an instrument delivered on time and on cost is the Moon Mineralogy
Mapper (M3), an imaging spectrometer designed to map the surface mineral composition of themoon for the Chandrayaan-1 Indian Space Research Organisation mission. This instrument was a
JPL in-house development managed by Mary White, the WFMOS project manager. The original
proposal cost was $23.6 M and final costs after delivery and preliminary integration on thespacecraft was $25.1 M, with the increase due to delays from the spacecraft provider. Completing M3 at cost required constant interaction with the spacecraft provider, detailed assessment ofdevelopment progress against schedule, and implementing design changes in a timely manner.
JPL involvement with ground instruments includes the Keck Interferometer and the M1 Control System (M1CS) for the Thirty Meter Telescope (TMT), and the calibration system for Gemini Planet Imager (discussed in section Management Proposal Section 12.1). JPL is also support-

70

WIDE FIELD FIBER-FED OPTICAL
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ing TMT on the design and development of the Alignment and Phasing system, telescope systemengineering, the Planet Formation Instrument feasibility study, and the telescope optics.
WFMOS aligns with JPL’s Astronomy and Physics Directorate (APD) strategic plan for increased development of ground and space-based optical instruments and JPL is committed to its
success. Ground instruments allow the development of complex systems with technologies that
can also be applied to space-flight instruments. JPL is a large, diverse organization accustomed
to handling complex, multi-disciplinary challenges and can provide critical resources in various
forms to support WFMOS. The support from the APD has been evident through the study phasethrough communication on the study progress and in the coordination and support in the proposal
review and preparation process. In addition to management support, the JPL infrastructure provides mechanisms and resources for cost and schedule management, document organization andstorage, and configuration control. This project has the advantage of selecting the desired tools
and functions without having to meet the more stringent practices of a flight project.
The WFMOS organization chart is shown in Figure 7.0-1.
GeminiDoug Simons, Director
Principal Investigator
Advisory Board

Joe Jensen, Head of Instrumentation

Richard Ellis, Caltech

Jakob Van Zyl, JPL
Andrew Lange, Caltech
Science Team

Ian Robson, UKATC
Project Management

George Efstathiou, IoA

Dark Energy: John Peacock, UKATC

Mary White PM

Larry Ramsey, PSU

Galactic Archaelology: Mike Irwin, IoA

Robin Bruno, DPM

Auxiliary Science: Ofer Lahav, UCL

Mike Seiffert, PS
GeminiScot Kleinman,
Instrument Program
Scientist

JPL

System Engineer
Justin Maki
JPL

Fiber System

Detector System

SpectrographMetrology System

System Software

Antonio Cesar de Oliveira

Roger Smith

Ian ParryPeter Doel

Steven Beard

Prime Focus Instrument&

LNA

CaltechIoA

UCL

Data Reduction

Positioner Assembly

Andy Vick

David Braun

UKATC

Acquisition and Guide
Lewis Roberts

U.S.
UK

JPL

BrazilPositioner Software
Karl Reichard

External to Project

PSU

Science teamFigure 7.0-1: WFMOS Organization Chart

The management team consists of the PI, Project Manager, Deputy Project Manager, Project
Scientist, and System Engineer.
• The role of the PI is to lead the development of the science case and to act as the final
authority to confirm that the instrument design and build meets the science goals. This
includes coordination of the full science team to ensure the goals of the DETF are met
and interaction with the technical team to ensure the instrument development is meeting
cost, schedule, and performance commitments. This interaction with the technical team is71

WIDE FIELD FIBER-FED OPTICAL
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SUMMARY REPORT

•

•

•

primarily with, but not limited to, the project manager, project scientist, and instrument
system engineer.
The project manager is responsible for delivering an instrument that meets therequirements on schedule and at cost. Responsibilities include communication with
Gemini and JPL program management, coordination of the technical team, and managing
resources. In this case, the instrument has a firm cost cap, so a critical feature of the
program manager’s role is to anticipate any cost overruns and, if necessary, implement
descopes in a timely manner. These actions must be communicated clearly to the PI andto Gemini if they affect any requirements from the OCDD Level 1 Requirements or
Instrument Level.
The project scientist’s role is to provide the modeling and assessment necessary toensure the instrument capability is consistent with the science objectives derived in the
OCDD. In addition, he fulfills the role of providing a strong tie between the science andengineering team as the design and build of the instrument progresses. In this role, he is a
key player in assessing trades that affect instrument performance and determining andquantifying the impact on the science objectives.
The system engineer reports to the project manager, leads the technical architecture of
the instrument and makes sure the subsystem development is balanced with respect totechnical risk and complexity. In addition, the system engineer is responsible for risk
management. All of this requires a comprehensive and in-depth knowledge of the
instrument development and interaction between subsystems so that system-level tradesin design or approach can be made effectively. The formalization of these decisions isthrough the management of the Level 2 and Level 3 requirements. In addition, the systemengineer is responsible working with the Subaru technical representative to obtain thetelescope ICDs and requirements. Finally, the system engineer manages and enforces
configuration control over all governing documents.

7.2

Instrument Cost
The key elements in cost management are accurate cost estimates, detailed cost tracking, andtimely actions for cost control. The total estimated cost for the WFMOS instrument is $68.5M,
which includes a baseline cost of $58.3M and $10.2M in reserves (20%). The reserves do notinclude the Gemini-held reserves of 15%. Table 7.2-1 shows the costs by partner and Table 7.2-2
provides the costs by top-level WBS. The partner contracts held with Gemini include 15% reserves, which are included in the total cost column in table 7.2-1. Reserves above 15% are included in the JPL total cost to allow management of reserves across the project. The process used
to estimate necessary reserves is presented in section 5.2 of the Management and Cost Volume
and in section 7.5 below. Half of the instrument cost should meet the Gemini partner shares requirement as shown in Table 7.2-3. The project partners are comprised of the U.S., UK, and Brazil; half of the project costs are distributed to the partners as shown in Table 7.2-3.

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Table 7.2-1: WFMOS Costs by Partner
Partner
JPL
Cambridge
PSU
Caltech
LNA
UK-ATC
UCL
Subaru Provided Elements
Baseline Total
Reserves
Total w/ Reserves

Baseline Amount, $K
27,237
10,383
563
2,578
3,550
4,355
1,639
8,000
58,305
10,222
68,527

Total Cost w/ Distributed Reserves, $K
34,385
11,940
647
2,578
4,083
5,009
1,885
8,000

68,527

Table 7.2-2: Cost by Top-Level WBS
Cost by Year, $K
WBS Element

Institution

yr 1

yr 2

yr 3

yr 4

Total Reserves, Reserves, Cost w/
yr 5 Cost, $K
Reserves,%$K$K

1.0 Science

Caltech/JPL 317

210

215

218

221

1,181

20.0

236

1,418

2.0 Management

JPL

895

1,273

1,310

1,212 904

5,594

10.0

559

6,153

3.0 System Engineering JPL

485

528

530

509

2,328

17.3

404

2,732

4.0 Instrument System

7,791

11,702 11,068 4,566 2,284 37,411 24.1

9,022

45,534276

4.1 SpectrographCambridge 1,847

3,278

3,698

959

601

10,383 20.0

2,077

12,4594.2 Positioner

JPL/PSU

2,490

3,410

1,765

288

324

8,276

23.9

1,453

9,729

4.3 Detector System

Caltech/JPL 729

1,337

982

147

183

3,379

28.3

870

4,248

4.4 Fiber System

LNA

1,231

952

625

441

302

3,550

20.0

710

4,260

JPL

701

1,628

2,871

2,157 561

7,918

26.0

2,120

10,038

UKATC

391

400

692

526

257

2,266

25.0

567

2,833

403

697

435

48

56

1,639

20.0

328

1,967

0

0

284

794

624

1,701

28.3

481

2,182

318

555

839

230

148

2,089

20.0

418

2,507

1,655

1,655

1,655

1,655 1,379 8,000

0.0

0

8,000

10,222

68,527

4.5 Prime Focus
Instrument
4.6 System Software

4.8 Metrology/Calibration
UCL
system
5.0 Instrument Integration
JPL
and Test
6.0 Data Systems
UKATC
7.0 Subaru Provided
GeminiElements
TOTAL

11,462 15,923 15,901 9,184 5,836 58,305 20.3

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Table 7.2-3: Partner Shares and WFMOS Cost Distribution
Country
United States
United Kingdom
CanadaBrazilArgentina
Australia
Total

Required, %
46.80
28.15
13.60
4.45
4.65
2.36
100.00WFMOS, %
58.94
35.45
0.00
5.60
0.00
0.00
100.00Details of the cost estimate are provided in the Management and Cost Volume. The cost estimates obtained are bottoms up estimates from the doing organizations supplemented with costanalogies for similar instruments or subsystems. Cost guidelines were developed and distributed
to ensure consistency and completeness in the estimates. Vendor quotes have been obtained for
major procurements. Each estimate has been through the originating institution’s financial officeso that all rates and overheads are accurate. A series of informal and formal reviews were conducted by JPL, STFC and an external review board during the development of the proposal. The
reserves for each subsystem were estimated based on design complexity, design heritage, teamexperience, development effort, fidelity of vendor quotes, and overall risk in delivery.
Cost tracking and management will be critical on this project. The PM is responsible formanaging and reporting the project schedule and costs with support from the Business Management Team. Although Earned Value Management (EVM) is not required for the WFMOS project, the fundamental principles of EVM will be applied. The cost and schedule management will
be implemented and controlled using JPL's business management system, which provides an integrated project management system for defining work, organizing the work, assigning responsibility, developing schedules, authorizing work, establishing budgets, measuring performance,
analyzing variances, controlling costs, and summarizing the technical, schedule, and cost information required for effective control of a project.
For tracking, monthly financial reports from each partner will be submitted and then consolidated for management review. To minimize overhead, each partner will submit their standard
institutional financial report, and the Project Resource Analyst (PRA) will assemble a projectlevel report with key features used by the Project Manager for review. An accurate understanding of the design status with respect to the schedule is critical for assessing the cost status. The
status of major milestones with expected costs is a good global indicator of cost trends.
In addition, a project lien list will be developed and maintained to track cost changes and potential increases. The lien list contains items that may require additional funds, the decision datefor execution, and the steps necessary to finalize the decision. This lien list is used by the management team as a means to assess cost trends and they will approve liens based on cost, instrument performance, technical risk, and schedule risk.
Cost management is critical on this project, and the management approach reflects that by assigning clear responsibilities and allocating sufficient time for cost, schedule, and risk analysis.
Any cost trends that indicate likely overruns and the realization of costs incurred through executed liens will be managed through the implementation of descopes and allocation of reserves.
Decisions made to control costs without impacting instrument performance typically come in oneof several forms: 1) reduced expected performance while still meeting the requirements (minimum performance), 2) elimination of design options that could have led to improved perform-

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ance, or 3) added risk. A list of descopes and trigger dates will be used to balance any cost overruns within the framework of balancing risk in the instrument. The most important feature of thisapproach is to implement the descope in a timely manner. Features maintained in the descope list
are:
• Descope options, with implementation approach• The resulting acceptable reduced performance or risk position.
• The schedule milestones related to the required decisions.
• The budget impact if the options are exercised at the identified milestones.
• The survey success impact.
• Reserve Management.
This strategy has been used successfully on other projects that have faced similar cost caps
such as M3, recently delivered, and the Nuclear Spectroscopic Telescope Array (NuSTAR) project, currently in Phase C/D. A formal process that may include a trade study to review the impact and final approval from the management team is required to exercise the descope. The advisory board may be involved if the change is significant and Gemini will be consulted on any
descope that affects a Level 1 or Level 2 requirement. The challenge is that the time to exercise adescope effectively is often early and there is a tendency to want to defer in order to maintain allfeatures of the original design. A descope list will be developed further at the start of the study,
but at this time two high-level descopes are to reduce total number of positioners and to use
standard spectrograph CCDs instead of custom CCDs. For the positioners, there are several execution dates: 1) start of positioner contract, 2) parts procurement, 3) positioner assembly, and 4)
bench assembly. At each stage, the cost savings decreases, but this allows the flexibility to maintain the baseline design as long as possible. We can also generate multiple trigger dates in the
procurement of custom CCDs. Similarly, cost savings decrease with time, but the option retains
the advantages of the system for as long as possible.
Both of these descope options described effectively increase the survey time. Science goals
are still achievable, but the instrument is less efficient.

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WIDE FIELD FIBER-FED OPTICAL
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Table 7.2-4: Descope Options
Item
Use COTS
CCDs instead ofcustom CCDs

Decision
Date
Project start,
parts order,
integration

Performance
Impact30%
reduction inwavelength
coverage orresolution,
depending on
surveyDecreased
observing
efficiency

Likelihood of
Performance
ImpactHigh

Other
Subsystems
AffectedSpectrograph

High

Spectrograph,fibers,detectors

$1.2M

Cost
Reduction$260K

Reduce numberof positioners
from 2,400 to
1800

Positioner
CDR, parts
order, start ofpositioner
integration

Reduce numberof spectrographsfrom 3 to 2

Spectrograph
CDR

Decreased
observing
efficiency

High

Positioners,
fibers,detectors$2.5M

Reduce thenumber of
metrology
cameras in JPL
I&T

Metrology
CDR, JPL I&T

Decrease inthroughput

Low

Software,thermal system

$490K

7.3

Comment

Coupled w/
reducing
spectographs to
2. Costs onlyreflect reduction
number of
positioners.
Additional savingsin integration and
test
Coupled w/
reducing numberof postioners to1800.
After early testing,it may be possible
to go to one 4K x
4K camera.
Additional error inmeasuring
postioner
locations resultsin loss ofthroughput

Schedule
The baseline schedule is 58 months as shown in Figure 7.3-1. The baseline costs are developed against the 58-month schedule; project reserves are sufficient to cover development andbuild risk as well continuing workforce if schedule reserves, estimated at 5 months, are used. If
the schedule reserves are used fully, the total project timeline is 63 months for a final delivery
date of December 13, 2014. This schedule shows the critical path is through the fiber system.
However there are two major paths leading to instrument integration; this schedule shows these
two major tested subsystems delivering to the telescope within two months of each other.
Some of the major milestones are:•Kickoff Meeting.
•
Project System Requirements Review (SRR).
•
Project and subsystem Preliminary Design Reviews (PDR).
•
Project and subsystem Critical Design Reviews (CDR).
•
Integration and Test Readiness Review.
•
Delivery of test slit blocks with fibers to the spectrograph.

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WIDE FIELD FIBER-FED OPTICAL
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Positioner EM completion.
•
Detector delivery to JPL for coating and delta doping.
•
Software delivery to spectrograph and positioner/PFI I&T.
•
Spectrograph/detector system ship to Subaru.
•
All final subsystem hardware deliveries to I&T.
•
Software and Data Pipeline delivery to Subaru.
•
Positioner/PFI ship to Subaru.
•
Completion of ATP.
Successful completion of major reviews according to the standards associated with SRR,
PDR, CDR, and PSR are good indicators of overall progress as well as the completion of significant tests and delivery of hardware assemblies. The detailed set of milestones along with clear
definition of success will be developed further at the start of the study.
Schedule tracking and management are important to this project as they are directly related tocost. Any schedule slip in one area can result in not only additional effort in that area, but holding personnel longer in other subsystems if it also impacts the final delivery date. A similar approach to the cost tracking is taken, collecting schedule status from all the partners on a monthly
basis and using a single integrated project schedule to assess the trends.
Options for schedule recovery are also actions that typically do not impact performance, but
may increase development risk. These can be items such as increasing workforce to acceleratetasks, replacing analysis with test, eliminating lower-level tests and capturing those objectives athigher level tests, staggered deliveries, and descopes. Depending on project costs, expedite feesto vendors may be used.
•

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WIDE FIELD FIBER-FED OPTICAL
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SUMMARY REPORTWFMOS Top Level Schedule
2009
F M A M J

J2010

A S O N D J F M A M J

10/1
MILESTONES2011
A S O N D

J F M A M J

J

8/2

PHASE B
1/15
SRR/Kickoff

J

PDR
3/2

CDR
7/8

Del Slit Block
7/15

10/5

Detector System

Del Slit
Blk Dsn
12/8

SRR
11/2

Dsn
Rvw LL
PDR
4/7
6/29

Prelim

PDR
CDR
4/14
7/12

InstalCorrector
Lens
4/11

PDR
1/12
DetailA&G Cameras
Ready for I&T
3/14

MGSE
Avail
3/29
7/15
Slit Block &Test Cables

Sys 1
11/9

Sys 3 Assy

6/5

Align

Del2
6/5

Slack

Test
& Int

SpectrographShip by Boat

PSR
7/1

Del3
11/6

Camera Camera Camera Camera
1
2
3
4
4/26
7/10
9/19

2/14

4/15 Del SCU

Ver2Recipes
4/4

Ver3
Ship HW
Recipes to JPL
7/29
11/5
Del Sys SW
3/13
Ver.
GeminiControl Accept
Setup Loop
Test ShipPDR CDR
5/31 7/27

Cable A
Rec Prelim
Sys Software

JPL Milestones

8/1 PH E

7/19 Del PFI Assembly

11/6

12/29

2/19

10/4
CamerasIntegration @ Subaru

Milestone

Rec
Fit
Check
5/6

CDR
6/8

Instrument Integration@ JPL

A S O N D

7/2 Install Fiducial Illuminator Sys
7/17 Install Thermal Control Sys
6/25 Install Power Distribution Sys

Del Optics,
Mech.
Parts& Electr. Sys 2
5/9

CDR
4/5

PDR
3/31

J

5/15 Del PFI

Fiber Cable B

CDR
6/7
Cal

2014
A S O N D J F M A M J

12/12
PSR/SIRDel Fiber Sys
Cable A
12/29

Detector Sys Del1
11/9

PDR
7/1

Data Reduction

System Software

CDR
2/23

CDR
10/28

Design Report
6/25

Metrology System
Spectral Cal. Unit

LL Optics

PDR
3/15

J

PHASE D

PFI/Positioner Integration @ JPL

Spectrograph

J F M A M J

10/15

CDR
11/10

Fiber Cable A

2013
A S O N D

Del Bench
6/20

PDR
1/28

LL

J

9/12
Integration & TRR

CDR Positioner
3/28 EM Sys
4/11

Prime Focus
Instrument
Fiber SystemF M A M J

8/1
Project CDR

PDR
5/17

Positioner

J

PHASE C

7/1
Project PDR

INSTRUMENT2012
A S O N D

Delivery Milestones

Figure 7.3-1: Top Level Schedule

78

subsystem CDR

SpectographArrivesubsystem PDR

Del
WFMOS
Instrument
7/11

9/5

Operations

Critical Path

LL Long Lead

WIDE FIELD FIBER-FED OPTICAL
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SUMMARY REPORT

7.4

Risk Management
Risk Management is used to balance the project risk position across all project resource areas(instrument success, technical reserves, schedule and cost reserves), controlling the distribution
and magnitude of the identified risks against the cost constraints while obtaining the best possible confidence in achieving high science return. Accurate assessment of the risks and taking action in a timely manner is critical to managing the technical delivery and cost of this instrument.
The risk management strategy consists of risk identification, risk assessment, and risk mitigation. Risk identification defines the risk, likelihood of occurrence, and impact of the occurrence
and is captured in a standard format as shown in Figure 7.4-1. All elements of the project, including reviewers, are involved in risk identification, and the risk list is managed by the system engineer. The risks and their ranking shown here represent the knowledge at this stage of the design.
Risk items are reviewed monthly to track change in status as the instrument development proceeds and to identify whether any mitigation strategies need to be exercised. The top high-levelrisks for the WFMOS instrument are shown on this chart.
Risk Assessment Criteria for Consequence and Likelihood
5
Risk Type (either or both may apply to each risk)
Operational Risk
Implementation Risk

LIKELIHOOD

4

Consequence of Occurrence

2

3

4, 5

2

3

1

3

4

Level
5
4
3
2
1

6, 7

1
1

2

5

CONSEQUENCE

Implement newprocess(es) or changebaseline plan(s)Criticality

Implementation Risk Level Definition
Cannot complete with current budget and resources
Consume all contingency, budget or schedule, >90%
Significant reduction in contingency or schedule slack, >50%
Small reduction in contingency or schedule slack, >10%
Minimal reduction in contingency or schedule slack, >1%

Likelihood of Occurrence
Level
5
4
3
2
1

Likelihood
Very HighHigh
Moderate
Low
Very Low

Aggressively manage;
consider alternative
process
MonitorFigure 7.4-1: Summary of Risk Assessment and Management.

79

Level Definition (*)
almost certain > 70%
More likely than not > 50%
Significant likelihood >30%
Unlikely > 1%
Very unlikely <1%

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Table 7.4-1: Summary of Risk Assessment and Management
ID
1
2
3
4
5
6
7

Item
Custom detectors don't meet
performance requirement
schedule risk on the spectrographgratings
The square wave drive approach
does not adequately drive thepositioner motors.
Positioner resolution too high or
velocity versus position and/or step
size too erratic.
Telescope environment not as
expected and affects instrument
performance
Positioner contractor unable toperform
Spectrograph Aspherics difficult tomanufacture

Likelihood

Consequence

2

4

3

2

2

3

2

2

Mitigation
Change to a 6k×4k format from Hamamatsu with
15µm pixels
Start testing w/ one grating or early low-spec grating.
Consider Kaiser Optics as a second supplier
Work w/ NST to modify motor design. Return to thesine wave drive and switching network - size growth
could exceed envelope
Enhanced positioner learning algortihm, Designmodifications after EM characterization2

2

Early (design phase) test and/or characterization oftelescope environment and interfaces

1

2

1

2

JPL support to New Scale w/ engineering, facilities,
other resources.
consider multiple vendors, manage aggressively, blockoff out of spec areas

Table 7.4-1 presents the current risks. The probability of occurrence is indicated by the likelihood rating. A short description of the risk and the mitigation plan is provided below.
1. The square wave drive approach does not adequately drive the positioner motors. Proper
packaging design and techniques are necessary to make the 2.5-volt signals from the
FPGA compatible with the 120-volt square wave out of the drivers. Testing will be
performed to ensure noise and crosstalk issues are addressed. An electronics design thatutilizes sine wave drive signals, increases isolation, and increases packaging volume ofthe electronics. If the noise and crosstalk cannot be minimized, the existing motor
electronics design from NST may be repackaged as a third option.
2. Positioner Resolution too high or velocity vs. position and/or step size erratic. The
current time allocation for repositioning fibers to a new target within 5 µm is 80 s. The
completed positioner bench test indicates reposition can actually be done in 40 sec;
however, the unknown of the consistency among positioners as well as any unpredicted
errors when integrating the full system prevents retiring this risk completely. The current
risk assessment represents growth in timeline. Modest design modifications and/or
enhanced positioner software algorithms will be employed if needed after early testingwith the positioner module EM.
3. Custom detectors do not meet performance requirement. The baseline design has custom
detectors from LBL with custom coatings from JPL. The CCD and coating designs are
not risky since they will be assembled from proven components and processes, however,this is dependent on a single resource at LBL. An alternate supplier (Hamamatsu) has
been identified in case LBL withdraws for any reason. The penalty will be higher costand lower sensitivity in the blue and extreme red.
4. Telescope environment not as expected and affects instrument performance. Additional
information on the thermal and dynamic environment of the telescope may reveal thestability of the coupled system of the telescope and instrument does not meet
performance requirements. Mitigation is early acquisition of additional data on the80

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telescope environment and interfaces in critical locations and/or simple measurements tocharacterize the local environments.
5. Some spectrograph gratings do not meet the delivery date. This schedule assumesdelivery of all gratings is in time for integration and test. Sufficient testing and early
integration can be done with a low-spec grating and/or a single accepted grating.
6. Positioner Contractor unable to Perform. This acknowledges a single source for a criticaldeliverable. Mitigate with close monitoring and supplement with JPL resources if needed.
This approach has been used often in the past with good success.
7. Spectrograph aspherics difficult to manufacture. This risk is listed because only verbal
quotes have been obtained from the vendors. It is expected to be addressed quickly by
seeking alternate vendors as well as active involvement and management of the selected
vendor.
Risk decisions are supported by analyses and recommendations from the project team, but
are ultimately made by the Project Manager in the same manner that all cost, schedule and performance impact decisions are made. If the scope of any change impacts Level 1 or 2 requirements, the management team, the advisory board, and Gemini are part of the decision process.
The descope list provides the leverage for accepting reasonable risk and maintaining confidence that a meaningful instrument can be developed within the fixed budget even if many of theidentified risks are actualized. The plan identifies functional or performance capabilities, which
may be dropped at some later point of the project, allowing cost savings to be applied to the area
where risks happened. It also describes where risks may be added (through dropping a test, orother activity with resulting cost avoidance) acceptably in order to cover unplanned costs in otherareas. As with the risk list, the descope list will evolve through the instrument development andneeds to be continuously assessed and updated.
7.5

Reserves
The total project reserves required is estimated at 20%; this 20% excludes the $8M allocatedfor the Subaru provided elements and does not factor in the Gemini-held contingency. The reserves are established by assessing risks associated with each WBS element (Management andCost Volume, section 5.2). Higher reserves are allocated to areas that: 1) require more development effort and therefore have more uncertainty in meeting performance cost and schedule, 2)
show risk of cost growth on procurements, 3) have a higher level of complexity in design orschedule coordination, and 4) have two or more identified significant risks at the project start.
The reserves do not explicitly include an allocation for exchange rate uncertainty and instead a
conservative exchange rate has been used to absorb rate fluctuations.
Reserves above the partners’ contract-required 15% are held at JPL and managed at the project level to provide maximum flexibility in risk management. Allocation of reserves is a processthat references the specific risk item and balances resources according to the summary of risks,schedule, and remaining resources. The project risks and reserves are reviewed and updated
monthly with the project team and Gemini and quarterly with the advisory board. A clear assessment and justification for the release of reserves is developed with the subsystem leads toassess the impact across the instrument system. The final decision is made by the PM with concurrence from the PI. The advisory board will be notified of the potential for significant changes
as they arise. A minimum level of reserves for each project phase is defined and managed to ensure sufficient funding is available through delivery. This is discussed further in the Management
and Cost Volume.
81

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8.0

References

Goodsell, Stephen J.; Blanken, Maarten F.; Corradi, Romano; Dee, Kevin M.; and Jolley, P. D.,“AutoFib2: small fibers,” Proc. SPIE 4841, 1367–1375, 2003.
Guzzo, Luigi et al., Nature 451, 541–544, 2008.
Parry, Ian R.; Lewis, Ian J.; Sharples, Ray M.; Dodsworth, G. N.; Webster, J.; Gellatly, David
W.; Jones, Lewis R.; and Watson, Frederick G., “Autofib--2: an automated fiber positioner
for the prime focus of the William Herschel Telescope,” Proc. SPIE 2198, 125–133, 1994.
White, M., Song, Y.-S., and Percival, W.J., 2008. arXiv:0810.1518.

82

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9.0

Acronyms and Abbreviations

1D
2dF
2dFDR
2dFGRS
3D
A&G
A/R
AAO
AAT
ACSIS
ACTDP
AD
ADC
ADU
AES
AF2
AG
AGC
AIV
ALMA
AMNH
ANA
AO
APD
APG
API
APS
AR
ARC
ARGOS
ARL
ARR
ATC
ATP
AURA
BAO
BBAR
bHROS
BIT
BOM
BOSS
BRI
CAD
Caltech
CAS
CCB
CCD
CCW
CDD
CDR
CELT
CFHT

One-Dimensional
The 2 Degree Field instrument
2dF Data Reduction
The 2dF Galaxy Redshift Survey
Three-Dimensional
Acquisition and Guiding
Antireflection
Anglo-Australian Observatory
Anglo-Australian Telescope
Auto-Correlation Spectrometer and Imaging System
Advanced Cryocooler Technology Development Programs
Applicable Document
Atmospheric Dispersion Compensator
Analog-to-Digital Unit
Auger Electron Spectroscopy
Auto-fib2
Autoguider
Acquisition and Guidance Camera
Assembly, Integration, and Verification
Atacama Large Millimeter Array
American Museum of Natural History
Analysis workstation
Adaptive Optics
Astronomy and Physics Directorate (JPL)
Annealed Pyrolytic Graphite
Application Programming Interface
Alignment and Phasing System
Antireflection
Astronomical Research Cameras
Abundance and Radial Velocity Galactic Origins and Structure Survey
Applied Research Laboratory
Acceptance Readiness Review
Astronomy Technology Centre
Acceptance Test Plan
Association of Universities for Research in Astronomy
Baryonic Acoustic Oscillations
Broadband AR
Bench-mounted High Resolution Optical Spectrograph (Gemini)

Built in Test

Bill of Material
Baryon Oscillation Spectroscopic Survey
Standard photometric bands (Johnson)
Computer-Aided Design
California Institute of Technology
Central Alarm System
Change Control Board
Charge-Coupled Device
Counter-Clockwise
Concept Design Documentation
Critical Design Review
California Extremely Large Telescope
Canada-France-Hawaii Telescope
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CGS4
CIB
CIRPASS
CIT
CL
CLD
CM
CMB
CMII
CMM
CMOS
CMP
CMS
CNC
COB
CoDR
COO
CosmoMC
COTS
CPU
CR
CRB
CSRP
CTE
CTIO
CTM
CV
CVD
CW
DAQtk
DAZLE
DBSP
DC
DCS
DE
DEC
DEHZ
DE-HZ
DEIMOS
DELZ
DETF
DHS
DID
DLD
DMA
DoD
DOE
DOF
DOORS
DPM
DR
DRR
DS
DSP

Cooled Grating Spectrometer 4
Configuration Implementation Board
Cambridge Infrared Panoramic Survey Spectrograph
California Institute of Technology
Center Line; also confidence level
Center for Life Detection
Configuration Management
Cosmic Microwave Background
Configuration Management II
Coordinate Measuring Machine
Complementary Metal Oxide Semiconductor
Configuration Management Plan
Configuration Management System
Computer Numerical Control
Chip on Board
Concept Design Review
Caltech Optical Observatories
A Markov-Chain Monte-Carlo software package
Commercial-Off-the-Shelf
Central Processing Unit
Change Request; also Cosmic Ray
Configuration Review Board
Center for Space Research Programs
Charge Transfer Energy
Cerro Tololo Inter-American Observatory
Contract Technical Manager
Curriculum Vitae
Chemical Vapor Deposition
Clockwise
Data Acquisition Toolkit
Dark Age Z Lyman-alpha Explorer
Double Spectrograph at Hale 200" Cass focus
Direct Current
Detector Control Software
Dark Energy
Declination
Dark Energy, High z / High Redshift Dark Energy
Dark Energy, High z / High Redshift Dark Energy
Deep Imaging Multi-Object Spectrograph (Keck)
Dark Energy, Low z / Low Redshift Dark Energy
Dark Energy Task Force
Data Handling System
Data Item Description
Digital Logic Decoder
Direct Memory Access
Department of Defense
Department of Energy
Degrees of Freedom
Dynamic Object Oriented Requirements System
Deputy Project Manager
Data Reduction
Design Release Review
Data Set
Digital Signal Processor
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E2V
ECN
ECoS
EGSE
ELT
EM
EMC
EMI
ESA
ESD
ESO
EU
FAR
FEA
FIS
FITS
FLAMES
FLCG
FMEA
FMOS
FOCCOS
FOV
FPC
FPGA
FPRD
FRD
FRS
FS
FSL
FTE
Fundep
FWHM
GA
GAHR
GA-HR
Gaia
GALEX
GALR
GA-LR
GCal
Gen2
G-ICD
GMOS
GOODS
GPI
GPMC
GPol
gri
GRS
GSA
GUI
H/W
HALT

e2v Technologies PLC, Essex, England
Enterprise Change Notice; also Engineering Change Notice
Eberly College of Science
Electrical Ground Support Equipment
Extremely Large TelescopeEngineering Model
Electromagnetic Compatibility
Electromagnetic Interference
European Space Agency
Electro-Static Discharge
European Southern Observatory
European Union
Final Acceptance Review
Finite Element Analysis
Fiducial Illumination System
Flexible Image Transport System
Fiber Large Array Multi Element Spectrograph
Field Lens Collimator Group
Failure Modes Effects and Analysis
Faint Multi-Object Spectrograph
Fiber Optical Cable and Connector System
Field of View
Fiber Positioner Controller
Field-Programmable Gate Array
Functional Performance Requirements Document
Focal Ratio Degradation
Fiber Routing System
Feasibility Study; also Frame Set
Frequency Standards Laboratory
Full-Time Equivalent
Fundação de Desenvolvimento da Pesquisa
Full Width at Half Maximum
Galactic Archaeology
Galactic Archaeology–High Resolution / High Resolution Galactic Archaeology
Galactic Archaeology–High Resolution / High Resolution Galactic Archaeology
An ESA-led astrometry space mission
Galaxy Evolution Explorer
Galactic Archaeology–Low Resolution / Low Resolution Galactic Archaeology
Galactic Archaeology–Low Resolution / Low Resolution Galactic Archaeology
Gemini Calibration Unit
.Generation 2.
(next generation Subaru Operations System Software (SOSS))
Gemini ICD
Gemini Multi-Object Spectrograph
Great Observatory Origins Deep Survey
Gemini Planet Imager
Governing Program Management Council
Gemini PolarimeterStandard photometric bands (SDSS)
Galaxy Redshift Survey
Gemini Science Archive
Graphical User Interface
HardwareHighly Accelerated Life Test
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HASSHETDEXHOMERHRHRNIRS
HSC
HTTP
I&T
ICD
ICM
ICS
IFS
IFU
IoA
IP
IPAC
IR
IRAF
IRMOS
IRR
IS
ISOCAM
ISW
IT
ITL
ITO
IVOA
JAC OT
JCMT
JFET
JHU
JPL
JWST
K.A.O.S.
KIDS
KPP
KVM-o-IP
LARS
LBL
LBNL
LDB
LED
LFC
LGS
LISA
LNA
LOE
LR
LRIS
LRU
M1CS
M3
MATLAB
MAXR

Highly Accelerated Stress Streening
Hobby Eberly Telescope Dark Energy ExperimentHigh Altitude Ozone Monitoring Experimental Rocket
High Resolution
High Resolution Near Infrared Spectrograph
HyperSuprimeCam
Hypertext Transfer Protocol
Integration and Test
Interface Control Document
Institute of Configuration Management
Instrument Control Software
Integral Field SpectrographIntegral Field Unit
Institute of Astronomy (Cambridge)
Internet Protocol
Infrared Processing and Analysis Center
Infrareda library of astronomical data reduction software
Infrared Multi-Object SpectrographIntegration Readiness Review
Investigation Scientist; also Information System
Infrared Space Observatory Camera
Integrated Sachs-Wolfe
Information Technology
Imaging Technology Laboratory, University of Arizona
Indium Tin Oxide
International Virtual Observatory Alliance
Joint Astronomy Center Observing Tool
James Clerk Maxwell Telescope
Junction Field Effect Transistor
Johns Hopkins University
Jet Propulsion Laboratory
James Webb Space Telescope
Kilo-Aperture Optical SpectrographKilo-Degree Survey
Key Performance Parameter
Keyboard, Video, Mouse over Internet Protocol
Laboratory for Reliable Software
Lawrence Berkeley Laboratory
Lawrence Berkeley National Laboratories
Logic Decoding Board
Light-Emitting Diode
The Palomar Large-Format Camera
Laser Guide Star
Laser Interferometer Space Antenna
Laboratório Nacional de Astrofisica
Level of Effort
Low Resolution
Low Resolution Imaging SpectrometerLine Replaceable Unit
M1 Control System
Moon Mineralogy Mapper
Programming language for technical computing from The MathWorks
The maximum resolution spectrograph mode86

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MBE
MCT
MDD
MDL
ME
Michelle
MIP
MIRI
MLS
MMR
MMT
MNRAS
MOE
MOP
MOS
MOSFET
MPO
MPS
MRO
MROI
MRR
MS
MT
MTBF
MTHR
MTO
MUSE
N&S
NA
NAOJ
NAS
NASA
NDEAA
NIR
NOAO
NSCAT
NSF
NST
OBCP
OBS
OCD
OCDD
OCS
OD
OFHC
OFS
OG
OIO
OLDP
OPE
ORB
ORR
OSL

Molecular Beam Epitaxy
Ministry of Science and Technology
Mechanical Design Description
Micro Devices Laboratory
Mechanical Engineer
Mid-Infrared Echelle SpectrographMaterial Inventory Plan
Mid-Infrared Instrument
Microwave Limb SounderMonthly Management Review
Multiple Mirror TelescopeMonthly Notices of the Royal Astronomical Society
Measure of EffectivenessMeasure of Performance
Multi-Object Spectrograph
Metal Oxide Semiconductor Field Effect Transistor
Multi-Fiber Push-on (a type of fiber optic connector)
Movement Planning Software
Mars Reconnaissance Orbiter
Magdalena Ridge Optical Interferometer
Make Release Review; also Manufacture Readiness Review
Mechanical Specification
Media-termination (a type of connector)
Mean Time Between Failures
Moderate to High Resolution Spectrograph
Motor Toggle OrderMulti-Object Spectroscopic Explorer (ESO VLT)
Nod and Shuffle
Not Applicable; also Numerical Aperture
National Astronomical Observatory of Japan
Network Attached Storage
National Aeronautics and Space Administration
Nondestructive Evaluation and Advanced Actuators
Near-Infrared
National Optical Astronomy Observatory
NASA Scatterometer
National Science Foundation
New Scale Technologies
Observation Control Processor
Operational Business Schedule
Operational Concept Design
Operational Concept Design Document; also Operational Concepts Definition
Document
Observatory Control System; also Observation Control Software
Optical Designer; also Outer Diameter
Oxygen Free High Conductivity
Fiber connector manufacturer
Output Gate
Oliveira Instrumentos Opticos
On-Line Data Processing
Observation Procedure Editor
Orbital Sciences
Operations Readiness Review
Optical Science Laboratory
87

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OSM
OWS
OzPoz
P2PP
PALM-LGSPanSTARRS
PASP
PBS
PC
PCA
PCB
PCD
PCI
PCO
PDMS
PDR
PDS
PECVD
PFC
PFI
PFU
PHSC
PI
PM
PMS
POP
POSM
PPARC
PRA
PRS
PS
PSA
PSE
PSF
PSR
PSU
PTFE
PTI
QA
QC
QE
QH
QTH
RA
RAVE
RD
Req
RFP
RM
RMS
ROE
ROI
ROM
ROMPS

Object Selection Mechanism
Observation Workstation
A fiber positioner unit for the Very Large Telescope (VLT)
Phase 2 Proposal Preparation
Palomar Laser Guide Star
Panaromic Survey Telescope and Rapid Response System
Publications of the Astronomical Society of the Pacific
Product Breakdown Structure
Personal Computer
Principal Component Analysis
Printed Circuit Board
Product Concept Design
A computer bus standard
Project Close Out
Project Database Management System
Preliminary Design Review
Power Distribution System
Plasma Enhanced CVD
Prime Focus Corrector
Prime Focus Instrument
Prime Focus Unit
PFU from the HyperSuprimeCam
Project Initiation; also Principal Investigator
Project Manager
Project Master Schedule
Program Operating Plan
Prismatic Object Selection Mechanism
Particle Physics and Astronomy Research Council
Project Resource Analyst
Pipeline Reduction System
Project Scientist
Project Schedule Analyst
Project System Engineer
Point Spread Function
Pre-Ship Review
Pennsylvania State University
Polytetrafluoroetheylene – Teflon
Palomar Testbed Interferometer
Quality Assurance
Quality Control
Quantum Efficiency
Quartz Halogen; also Quiet Hour
Quartz Tungsten Halogen
Right Ascension
Radial Velocity Experiment
Reset Drain; also Reference Document
Requirement
Request for Proposal
Risk Management
Root Mean SquareRoyal Observatory Edinburgh
Region of Interest
Rough Order of Magnitude
Robotically Manipulated Plug-Plate System
88

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RR
RSD
RSS
RTD
S/N
S/W
SAA
SAMP
SAN
SARR
SC
SCP
SCUBA
SCUBA2
SDSS
SE
SEGUE
SHARAD
SIA
S-ICD
SIM
SimCam
SIR-C
SIRR
SKA
SKADS
SNAP
SNR
SOAR
SOSS
SOW
SPEC
SPIE
SPIRE
SPLAT
SQR
SRB
SRE
SRF
SRL
SRR
SRS
SSA
ST7
STARS
STCS
STELESSTFC
SUBARU
SW
SWIFT
TBA
TBC
TBD

Requirements Review
Redshift Space Distortions
Root Sum Squared
Resistance Temperature Detector; also Real Time Display
Signal to Noise
SoftwareSpace Act Agreement
Simple Application Messaging Protocol
Storage Area Network
System Architecture and Requirements Review
Subprime Cam
Service Connection Point
Submillimeter Common-User Bolometer Array
Submillimeter Common-User Bolometer Array 2
Sloan Digital Sky Survey
Systems Engineering
Sloan Extension for Galactic Understanding and Exploration
Shallow Radar
Simple Image AccessSubaru ICD
Space Interferometry Mission(a working name for the next generation instrument interface)
Spaceborne Imaging Radar–C
System Integration Readiness Review
Square Kilometer Array
Square Kilometre Array Design Studies
SuperNova Acceleration Probe
Signal-to-Noise Ratio
Southern Astrophysical Research Telescope
Subaru Operations System Software
Statement of WorkSpecification
Society of Photo-Optical Instrumentation Engineers
Spectral and Photometric Imaging Receiver
Spectral Analysis Tool
Rotary Squiggle Motors
Strain Relief Box
Spectral Resolution Element
Sequence Request File
Significant Risk List
System Requirements Review
Software Requirements Specification
SuperCosmos Science Archive
Space Technology 7
Subaru Telescope Archive System
Subaru TCS
SOAR Telescope Echelle Spectrograph
Science and Technology Facilities Council
Subaru Telescope Archive System
SoftwareShort Wavelength Integral Field Spectrograph
To Be Announced
To Be Confirmed
To Be Defined; also To Be Decided
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TBR
TCP
TCS
TES
TFCALC
ThAr
TIKI
TIM
TMC
TMT
TPF
TPM
TUE
UARC
UCL
UCLES
UK
UK ATC
UKIRT
USA
USM
UUV
UV
UVES
V
V&V
VCP
VDD
VIPERSVISTA
VISTAIRCAM
VLT
VME
VNC
VO
VPH
VPHG
WBS
WCS
WCSE
WFC
WFCAM
WFMOS
WFS
WHT
WM
WMAP
WWW
WY
WYFFOSXML

To Be Reviewed
Transmission Control Protocol
Telescope Control System
Top-End Structure
Software package used for designing thin films
Thorium Argon
Telescope Interface Kit
Technical Interchange MeetingTechnical and Management Committee
Thirty-Meter Telescope
Terrestrial Planet Finder
Technical Performance Measure
Top-End Unit Exchange
University Affiliated Research Center
University College London
University College London Echelle SpectrographUnited Kingdom
United Kingdom Astronomy Technology Centre
United Kingdom Infrared Telescope
United States of America
Ultrasonic Motor
Unmanned Undersea Vehicle
Ultraviolet
UV-Visual Echelle Spectrograph
Volts
Verification and Validation
Verification and Commissioning Plan
Version Description Document
VIMOS Public Extragalactic Redshift Survey
Visible and Infrared Survey Telescope for Astronomy
VISTA Infrared Camera
Very Large Telescope
Versa Module Europa
Virtual Network Computing
Virtual Observatory
Volume Phased Hologram
Volume Phase Holographic Grating
Work Breakdown Structure
WFMOS Control SystemWFMOS WCS Electronics
Wide Field Corrector
Wide-field camera (IR camera for UKIRT)
Wide Field Fiber-Fed Optical Multi-Object Spectrometer
Wavefront Sensor
William Herschel Telescope
Work MonthWilkinson Microwave Anisotropy Probe
World Wide Web
Work YearAn intermediate dispersion spectrograph for the WHTExtensible Markup Language90