High Resolution Optical Spectrograph Feasibility Study ...pjm/QSpec/Instrument...

High Resolution Optical Spectrograph Feasibility Study Report

(HROS FSR)

Date: February 15, 2006 Document Number: TMT.INS.CDD.06.004.DRF01 Revision: Initial Release Contract No.: TMT.INS.CON.05.015.REL01 CDRL No.: HROS FSR

Prepared By: C. Froning S. Osterman M. Beasley S. Beland 02/15/2006 Date Reviewed By: Date Reviewed By: Date Approved By: Date Approved By: Date Approved By: Date

Center for Astrophysics & Space Astronomy University of Colorado, UCB 593

Boulder, Colorado 80309

REVISIONS

Letter ECO No. Description Check Approved Date Initial Release 2/15/06 Original Release THE UNIVERSITY OF COLORADO

Name Date At Boulder Drawn: C.S. Froning 01.18.06 The Center for Astrophysics and Space Astronomy Reviewed: High Resolution Optical Spectrograph Approved: Feasibility Study Report (HROS FSR)

Size Code Indent No. Document No. Rev

A HROS-FSR-01 V1.0 Scale: N/A

HROS-FSR-01 February 15, 2006 Center for Astrophysics & Space Astronomy Version 1.0

High Resolution Optical Spectrograph Feasibility Study Report University of Colorado at Boulder Page i

1. Introduction.............................................................................................................. 1

1.1 Overview............................................................................................................ 1 1.2 Scope.................................................................................................................. 3 1.3 Acronyms and Abbreviations.............................................................................. 3 1.4 Reference Documents ......................................................................................... 4

2. Scientific Motivation................................................................................................ 4 2.1 Science Case Overview ...................................................................................... 4 2.2 HROS Role in the TMT Science Program........................................................... 5

3. Instrument Requirements Flowdown ........................................................................ 7 3.1 Recommended Modifications to Instrument Requirements ................................. 9

4. Instrument Concept .................................................................................................. 9 4.1 Concept Overview .............................................................................................. 9 4.2 Optical Design.................................................................................................. 11

4.2.1 Design Overview ....................................................................................... 11 4.2.2 Major Optical Subsystems and Components............................................... 13 4.2.3 Performance............................................................................................... 34

4.3 Mechanical Design ........................................................................................... 37 4.3.1 Physical Location and Dimensions............................................................. 37 4.3.2 Enclosure................................................................................................... 38 4.3.3 Component Parts........................................................................................ 38 4.3.4 Power Requirements and Output ................................................................ 38

4.4 Electronics Design............................................................................................ 39 4.4.1 Mechanisms............................................................................................... 39 4.4.2 Detector Electronics................................................................................... 42

4.5 Software Design ............................................................................................... 42 4.5.1 Instrument Control Software ...................................................................... 42 4.5.2 Data Flow and Handling ............................................................................ 43 4.5.3 Observation Planning Tools ....................................................................... 43 4.5.4 Data Analysis Tools................................................................................... 43

4.6 Key and High Risk Components ....................................................................... 44 4.6.1 FIFU.......................................................................................................... 44 4.6.2 Dichroic Tree............................................................................................. 44 4.6.3 Cameras..................................................................................................... 45 4.6.4 Detectors.................................................................................................... 45

4.7 Instrument Maintenance.................................................................................... 46 4.7.1 Maintenance Requirements ........................................................................ 46 4.7.2 Spare Parts List .......................................................................................... 46 4.7.3 Instrument Downtime ................................................................................ 47

5. Performance Estimates ........................................................................................... 47


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5.1 Sensitivity Calculator........................................................................................ 47 5.2 Technical Viability of Science Programs .......................................................... 49

6. Trade Studies.......................................................................................................... 49 6.1 Foreoptics......................................................................................................... 50

6.1.1 Atmospheric Dispersion Compensation...................................................... 50 6.1.2 Acquisition and Guiding ............................................................................ 50 6.1.3 Calibration ................................................................................................. 50

6.2 Adaptive Optics................................................................................................ 50 6.2.1 SLGLAO ................................................................................................... 51 6.2.2 GLAO........................................................................................................ 51

6.3 Options for Image Slicing................................................................................. 52 6.4 Observing Modes and Efficiency ...................................................................... 53

6.4.1 Spatial Sampling........................................................................................ 53 6.4.2 Sensitivity.................................................................................................. 53 6.4.3 Simultaneous Wavelength Range ............................................................... 54 6.4.4 Spectral Resolution .................................................................................... 54

6.5 Deployment on Nasmyth Platform.................................................................... 54 6.6 Vacuum Chamber ............................................................................................. 55

7. Upgrade Options .................................................................................................... 55 7.1 J-Band Arm ...................................................................................................... 55 7.2 Multi-Object Capability .................................................................................... 55 7.3 AO Interface: SLGLAO to GLAO................................................................... 56 7.4 R>100,000 Spectroscopy .................................................................................. 56

8. Descope Options .................................................................................................... 56 8.1 R≤60000 Only .................................................................................................. 56 8.2 R=100,000 Through Non-Contiguous Wavelength Coverage............................ 57 8.3 Blue Cutoff at 340 nm ...................................................................................... 57 8.4 Single Object Spectroscopy Only...................................................................... 57 8.5 Elimination of Laser Frequency Comb.............................................................. 58 8.6 Elimination of Vacuum Enclosure for Optics Cavity......................................... 58

9. Cost Estimates........................................................................................................ 58 9.1 Default Design.................................................................................................. 58

9.1.1 Budget Narrative........................................................................................ 59 9.2 Upgrade Costs .................................................................................................. 60 9.3 Descope Cost Savings....................................................................................... 60

10. Instrument Development Program ........................................................................ 61 10.1 Work Breakdown Structure............................................................................. 61 10.2 Development Schedule ................................................................................... 62

11. Bibliography......................................................................................................... 63


High Resolution Optical Spectrograph Feasibility Study Report University of Colorado at Boulder Page iii

List of Figures

Figure 4-1: Schematic diagram of the CU-HROS Concept. Light from the TMT is fed into the CU-HROS instrument enclosure and reconstructed into a 0.1" pseudo-slit via a fiber optic IFU. The spectrum is divided into 32 spectral bins via a 31 element dichroic mirror array, feeding 32 narrow band, first order spectrographs with near constant resolution and efficiency.......................................................................... 10

Figure 4-2: One possible layout of the dichroic array. Light enters from the upper left, and either passes through to the upper red arm, or is reflected back and folded into the blue arm (shown in green). The feed points for the individual spectrograph benches are indicated by spheres; the benches would either be vertical, mounted above and below the dichroic planes (requiring an additional fold mirror for each channel), or with the grating in the dichroic plane, and the dispersion plane vertical. The arrangement shown is not optimized, and is subject to more detailed analysis. Overall size is roughly 5x3x0.75 m without the spectrograph benches. .................. 11

Figure 4-3: Block diagram showing the major components of the CU-HROS design for a single channel........................................................................................................ 13

Figure 4-4 An example set of ADC prisms for the 70 arcsecond FOV that would be sufficient for our needs. ......................................................................................... 14

Figure 4-5: Consecutive pulses emitted by a mode-locked laser (a) and the corresponding spectrum(b). (from Udem, et al., 2002).................................................................. 20

Figure 4-6: (a) Illustration of one possible remapping of the fiber optic IFU input to the pseudo-slit output. Tessellated hexagonal microlenses (used to minimize geometrical FRD) are used to concentrate light onto the individual fibers, which are then remmaped into a long, narrow pseudo-slit. Various remapping schemes, such as the spiral mapping shown here, and various raster mappings will be considered prior to implementation. (b) Distribution of light at the pseudo-slit for a 0.5 arc second Gaussian PSF. Spiral image dissection concentrates light towards the center of the pseudo-slit. .................................................................................................. 23

Figure 4-7: Example collimator based on the formalism of Korsch (1991). The collimated beam exits at the lower left side of the panel......................................... 24

Figure 4-8: Measured performance for Wide Field 3 dichroic mirror F606W (blue line) delivered to NASA. An earlier, rejected mirror is shown for comparison (dashed red). Figure from WF3 monthly status review presentation, Sept. 1 2005 (John MacKently presenting). ......................................................................................... 26

Figure 4-9: Measured performance for an existing 125x125 mm dichroic mirror; data provided by Barr Associates. ................................................................................. 26

Figure 4-10: Simulation of a typical AGN spectrum before and after passing through the dichroic tree, based on Barr theoretical data for our requested mirror set. .............. 27

Figure 4-11: Performance of filter stacks for each channel based on theoretical data provided by Barr Associates. As in the previous figure, predicted performance is


High Resolution Optical Spectrograph Feasibility Study Report University of Colorado at Boulder Page iv

reduced to match specification. Roll off at short wavelengths is induced principally by the short wavelength performance of the initial filter, which must have excellent performance over an extremely wide wavelength range. Also note that high frequency fluctuations in individual filters do not correlate, but instead tend to average out after 5 transmissions/reflections. ......................................................... 29

Figure 4-12: Example of a six-element double-gauss camera suitable for CU-HROS. Channel 17 (569 - 594 nm) uses two aspheric elements (first and last lenses). Note anamorphism in spectral direction. Lenses are on the order of 250 mm.................. 32

Figure 4-13: Diagram of the different mechanisms and their links ................................ 41 Figure 10-1: CU-HROS Development Schedule........................................................... 63


High Resolution Optical Spectrograph Feasibility Study Report University of Colorado at Boulder Page v

List of Tables

Table 1: Members of the Team for the CU-led HROS Feasibility Study .........................3 Table 2: CU/HROS Performance Requirements/Goals ...................................................9 Table 3: Requirements for the derotator........................................................................ 15 Table 4: Sample HROS Dichroic Edges and Corresponding Channel Widths ................ 29 Table 5: Pixel Illumination for CU-HROS and a Conventional Echelle on TMT .......... 33 Table 6: Spectral coverage for CU-HROS over the optical waveband........................... 36 Table 7: Estimated throughputs for CU-HROS............................................................. 36 Table 8: Power Requirements for CU-HROS ............................................................... 39 Table 9: Spare parts list for CU-HROS.......................................................................... 47 Table 10: Limiting Magnitude for 6-hour total exposures with 5 reads, on-chip binning

for lower resolution. .............................................................................................. 49 Table 11: Comparing fiber IFUs to Image Slicers......................................................... 53 Table 12: Cost Estimates for the Default CU-HROS Concept....................................... 59 Table 13: Upgrade Costs for CU-HROS....................................................................... 60 Table 14: Descope Cost Savings for CU-HROS ........................................................... 61 Table 15: Top Level WBS for CU-HROS Construction................................................ 62


High Resolution Optical Spectrograph Feasibility Study Report University of Colorado at Boulder Page 1

1. INTRODUCTION

1.1 OVERVIEW

This document presents the results of a feasibility study of a High Resolution Optical Spectrograph (HROS) for the Thirty Meter Telescope (TMT). The study was conducted by a team led by the Center for Astrophysics and Space Astronomy (CASA) at the University of Colorado at Boulder (CU Boulder) from 2005 May through 2006 April. Table 1 shows the members of the study team and their responsibilities on the study. It is our view that HROS will be a key instrument for TMT, providing precision observations of Galactic and extragalactic targets that other planned ground-based and space-based instruments will be unable to obtain. Unlike many instruments on TMT, HROS will also be able to obtain high-quality science data when observing conditions are less than ideal and/or when adaptive optics (AO) systems aren’t operational. As a result of this combination of powerful science acquisition capabilities and flexibility of use, we consider HROS an excellent choice for a first light instrument. We have designed our instrument concept to be simple and stable, capable of benefiting from improved seeing conditions and AO but not dependant on these, and optimized for R=100,000 spectroscopy, which we believe will take advantage of the unique capabilities of TMT in the next generation of ground-based astronomy. This report is organized as follows: Section 2 reviews the science case for HROS (presented in detail in the CU-HROS IOCDD) and presents the expected role HROS will play as part of the TMT instrument complement; Section 3 reiterates the flowdown of the science requirements to the instrument requirements (as developed in the CU-HROS IFPRD); Section 4 presents our complete optical, mechanical, electronics, and software design for the CU-HROS concept, as well as discussion of key and high risk components and how they will be addressed, instrument maintenance, and implementation of the instrument on TMT; Section 5 gives sensitivity and performance estimates for the instrument; Section 6 presents trade studies for major system components; Sections 7 and 8 outline upgrade and descope options for the instrument design and associated costs; Sections 9 and 10 give the estimated cost and schedule for instrument design and construction; and finally, Section 11 lists the bibliography.



Name Institution Relevant Experience Responsibility

Cynthia Froning CASA/CU

COS Project Scientist; NIC-FPS Program Manager; compact objects & accretion disks

Feasibility Study PI; Science lead

Steven Osterman CASA/CU

COS Instrument Scientist, MAXIM, JMEX instrument teams

Technical Team Lead

Matthew Beasley CASA/CU UV sound rocket design; HORUS, DESTINY teams

Opto-mechanical design

James Green CASA/CU COS PI; FUSE spectrograph PI; NIC-FPS PI

Technical Team

Stephane Beland CASA/CU

COS, NIC-FPS, NRAO scientific programmer; CFHT instrument specialist

Software, Mechanisms

William Cochran U. Texas Planet searches; solar systems; stellar rotation

Science team

Inese Ivans Carnegie Stellar abundances; halo studies; chemical evolution

Science team

Steven Penton CASA/CU Lyα forest; COS instrument team Science team

Matthew Shetrone U. Texas

Stellar abundances; dwarf galaxy nucleosynthesis; HET HRS ops

Science team

J. Michael Shull CASA/CU IGM; baryonic structure; ISM Science team

John Stocke CASA/CU IGM; Lyα forest; galaxy evolution; QSOs

Science team

Kim Venn U. Victoria Cosmochemistry; extra-galactic stellar abundances

Science team



BATC Boulder Astronomical instrumentation Optics and detectors consulting

Table 1: Members of the Team for the CU-led HROS Feasibility Study

1.2 SCOPE

The CU-HROS science instrument is in the conceptual design phase. This report presents the results of a feasibility study for HROS, demonstrating that a high resolution optical spectrograph can be designed and constructed for TMT that meets the science goals laid out in the TMT Science-Based Requirements Document (SRC). It does not, therefore, represent a complete design study for the implementation of HROS on TMT; such a study would be required as a follow-on to this program when and if HROS development proceeds. Instead, this study is intended to provide sufficient information to the TMT Project to evaluate feasibility, performance requirements, interface requirements, performance estimates, and costs for HROS. This document should be read in conjunction with the CU-HROS IOCDD and IFPRD, which develop the science case for HROS and the flowdown of instrument requirements from the science requirements. This document will continue to be updated in response to new information and requests from the TMT Project and the HROS Review Panel through the end of the feasibility study in 2006 April.

1.3 ACRONYMS AND ABBREVIATIONS

ADC Atmospheric Dispersion Corrector

AO Adaptive Optics

AR Anti-reflective

Arcsec or " Arc second

CASA Center for Astrophysics and Space Astronomy

CCD Charge-Coupled Device

CU University of Colorado

CU-HROS University of Colorado Conceptual Design for HROS

FIFU Fiber Optic Integral Field Unit

FWHM Full Width at Half Maximum

GLAO Ground Layer Adaptive Optics

HROS High Resolution Optical Spectrograph

ICC Instrument Control Computer

ICS Instrument Control Software



IDL Interactive Data Language

IRAF Image Reduction and Analysis Facility

LGS Laser Guide Star

NGS Natural Guide Star

SLGLAO Single Laser Ground Layer Adaptive Optics

S/N Signal to Noise

TCS Telescope Control System

TMT Thirty Meter Telescope

WFS Wave Front Sensor

1.4 REFERENCE DOCUMENTS

RF1 Science-Based Requirements Document v15 (TMT.PSC.DRD.05.001.REL15)

Thirty Meter Telescope Project, January 21, 2005

RF2 Detailed Science Case v.9.0 (TMT.PSC.PRE.05.081.DRF01)

RF3 Initial Operations Concepts Definition Document

(TMT.INS.DRD.05.005.DRF03), CASA/CU, February 15, 2006

RF4 Initial Functional and Performance Requirements Document (TMT.

SEN.SPE.05.001.DRF03), CASA/CU, February 15, 2006

2. SCIENTIFIC MOTIVATION

2.1 SCIENCE CASE OVERVIEW

In the CU-HROS Initial Operational Concept Definition Document (IOCDD), we developed the science case for HROS and the subsequent requirements on the instrument to meet its science goals. Starting with the science case for TMT developed by the Science Advisory Committee (SAC) and presented in Detailed Science Case (DSC), we examined the present state of astrophysical research and worked to develop a picture of how HROS will advance this research in the next generation of ground-based astronomy. In the IOCDD, we presented detailed science cases for three broad areas: studies of the intergalactic medium (IGM) from the epoch of reionization to the modern universe,



planet searches and characterization of extra-solar planetary systems, and stellar abundance studies and cosmochemistry of the oldest stars in the Milky Way and out into isolated galaxies in the Local Group. Some of the highlights of the science cases include:

• Observations of the linear matter power spectrum of the IGM to determine the underlying dark matter structure of the universe at different redshifts;

• Tracing the spatial distribution and chemical composition of material flowing into and out of galaxies to determine the interaction history between collapsed structures and the reservoir of gas from which they form;

• Extending observations and abundance analyses out into the most diffuse components of the IGM to probe the cloud/intercloud boundary of the cosmic web;

• Expanding planet searches to a broader range of stellar types and ages and out into different components of the Milky Way (such as the Galactic bulge), with a dramatic increase in the number of cool stars available for study for searches for Earth-mass planets;

• Providing follow-up capabilities to observe Kepler transit systems to find Earth-mass planets and detect planetary atmospheric signatures;

• Undertaking “near-field cosmology” of a large population of very metal poor stars in the Milky Way halo to trace the formation and nucleosynthetic histories of their Population III progenitors;

• Extending stellar spectroscopy out into the populations of primordial isolated dwarf irregular galaxies to test models of large structure formation in the universe.

For each science case, we presented sample observing programs to demonstrate how HROS would be operated to obtain the necessary observations and then developed a set of instrument requirements subsequently used to shape the instrument conceptual design. For reasons we discuss in the next section, we also endeavored to design HROS to be flexible for a broad array of science observations, beyond those outlined above, to ensure its capability as a powerful, general-purpose facility instrument for TMT.

2.2 HROS ROLE IN THE TMT SCIENCE PROGRAM

In the IOCDD, we discussed specific cases in which we believe HROS will be at the forefront of advancing physical and astronomical science in the next generation of ground-based observations. History has demonstrated, however, that attempts to predict the scientific frontier a decade or more in advance are fraught with uncertainty. The process of predicting future scientific observations is a valuable one, as it helps to focus the discussion of instrument and telescope development on the prime goal of obtaining quality science data, but it nevertheless remains true that a review of the original



predicted science goals for instruments and telescopes currently in use largely failed to anticipate the most dramatic and the most profound results these facilities have enabled. As a result, any evaluation of the need for HROS on TMT must also take into account its place as a general facility instrument and an evaluation of how necessary it is to the overall success of TMT. An examination of the current generation of high resolution optical spectrographs on today’s 8 –10 meter class telescopes is illuminating in this regard. These instruments, particularly HIRES on the Keck telescope and UVES on the VLT, are currently driving progress in a number of fields and have been responsible for some of the most exciting discovery science of the past decade including, most notably, the discovery of extra-solar planets. When coupled with large collecting area, these instruments provide high resolution spectroscopy (with observations typically obtained at R~50,000) of faint targets (down to V~18 at moderate S/N). This potent combination has resulted in cutting-edge science, including: a revolution in the field of observational cosmology through observation of the IGM and galaxy-IGM interactions at redshifts from the epoch of reionization to the modern universe; constraints on theories of fundamental physics through observations of cosmological evolution in the value of the fine structure constant; the discovery of >130 extrasolar planets and the burgeoning study of where and how planets form; abundance analyses and nucleosynthesis histories of the oldest stars in the Milky Way and in nearby dwarf spheroidal galaxies; among many others. These instruments have also worked in complement with ground- and space-based low resolution spectrographs and imagers to provide the detailed spectroscopy necessary to conduct chemical, kinematical, and physical analyses of astrophysical targets. HROS will continue this tradition of excellence into the next generation of observational astronomy. HROS will leverage the >9 factor increase in collecting area TMT will provide over current telescopes to enable observations of currently accessible targets at twice the spectral resolution and S/N and to achieve observations of targets too faint for high resolution spectroscopy on current telescopes. The CU-HROS concept will support R=100,000 spectroscopy of the full optical waveband in a single exposure at S/N=100 for objects as faint as V~20, and R=20,000 spectroscopy at S/N = 20 to V<22. As a result, it will enable both high precision observations of faint targets at a combination of spectral resolution and S/N greater than can be achieved today and discovery-type lower resolution and S/N observations on targets too faint for today’s telescopes. The optical and redshifted ultraviolet wavebands covered by HROS contain the greatest density of information on the physical and chemical conditions of astrophysical gases of any portion of the electromagnetic spectrum, and therefore remain essential areas of astrophysical observation, even in the age of advances in infrared astronomy.



Finally, it is worth emphasizing that no comparable instruments are being planned for other next-generation missions likely to overlap with TMT: as a high resolution optical spectrograph on the premier ground-based telescope of the next decade, HROS’ role would be unique, and it would act as an essential complement to other missions. The historical record indicates the powerful scientific results that are obtained with the combination of a high resolution optical spectrograph and a large telescope, and underscore the need for HROS on the TMT.

3. INSTRUMENT REQUIREMENTS FLOWDOWN

From the science case, we have developed a list of instrument requirements. The detailed flowdown from science requirements to instrument requirements is presented in the IOCDD and the IFPRD. A summary of the performance requirements and goals was presented in Section 3.3 of the IFPRD and is reproduced in Table 2 below. Briefly, our science case indicates that HROS must cover the full optical waveband and even, within reasonable cost and the availability of other TMT instruments, extend into the non-thermal infrared. The full waveband should be covered in one exposure (or at worst, two exposures). Both of these requirements are driven by the presence of spectral features of interest for single projects over the full optical band: maximizing wavelength coverage maximizes the number of redshifts that can be probed in the IGM along distant sightlines as well as maximizing the number of useful redshifted UV transitions available for individual redshifts. Stellar abundance analyses (and, to a lesser extent, planet searches around stars) also rely on elemental lines across the optical waveband, and observing efficiency, particularly for faint stars, is maximized by obtaining all the transitions of interest at once. For IGM observations, planet searches, and stellar abundance analyses, a spectral resolution of at least R=50,000 must be attainable. This is sufficient resolution for radial velocity planet searches and abundance analyses of metal poor stars. A spectral resolution of up to R=100,000 should be aggressively pursued, however. At the higher resolution, weaker features in the IGM, the ISM, and in stars are resolved, allowing for probes of the diffuse IGM, kinematics in denser IGM structures, observations of relatively weak metal lines in stars, and spectral separation of lines in more metal-rich targets. HROS must maintain its aperture advantage over current instruments. Indeed, the HROS design must endeavor to maximize its throughput at all wavelengths over the optical waveband. Options for improving the seeing disk of the target and thereby improving the sensitivity of observations and reducing their exposure times should be pursued. These include the use of tip-tilt guiding correction to remove some atmospheric disturbance and the effects of wind shake and the ability of the instrument to interface with the TMT adaptive optics system at the time of HROS installation or at a later date. To facilitate



high precision velocity studies, such as those used for planet searches, HROS must be a stable instrument that achieves velocity precisions after calibration of <3 m/s and ideally <1 m/s. HROS must support stable, correctable observations both on a night to night level and over multi-year timescales. The field of view should be at least 70" to allow for the acquisition of natural guide stars for fast guiding and, potentially, adaptive optics corrections. Multi-object spectroscopy is useful for many programs, such as spectroscopy of stars in dwarf galaxies and lensed QSOs, but the relatively small size of the HROS field of view limits the breadth of applicability for multiplexing.

Specification Requirement Goal

Wavelength Coverage

310 – 1100 nm in a single, or at most two, observations. This may fall to 340 – 1100 nm depending on TMT optical coatings

310 – 130 nm in a single observation

Instantaneous Wavelength Coverage

One half of the full instrument band pass.

Full instrument band pass

Spatial Sampling <0.2" per lenslet/image slice at focal plane

0.1" per lenslet/image slice at focal plane

Spatial Resolution

<0.2" FWHM at detector <0.2" FWHM at detector

Field of View 10" 70"

Effective Slit Length

1"x5" 7 discretely positionable 1"x1" IFUs

Sensitivity Must maintain 30m advantage

Instrument efficiency shall be greater than 15% over 80% of band pass, and no less than 10% at any wavelength, including coaddition of distributed data

Support Calibrated Spectroscopy

Yes Yes

Spectral Resolution

R=50,000 R=100,000

Spectral Stability Line center shall be stable to 3 m/s for the duration of an observation

Line center shall be stable to 3 m/s for the duration of an observation

Spectral Wavelength calibration must Wavelength calibration must be



Repeatability be capable of providing centroid knowledge to <3 m/s between observations

capable of providing centroid knowledge to <1 m/s between observations

Table 2: CU/HROS Performance Requirements/Goals

3.1 RECOMMENDED MODIFICATIONS TO INSTRUMENT REQUIREMENTS

The required wavelength coverage should be altered to reflect the choice of mirror coatings for TMT as soon as that choice is made. We have altered several of the requirements to reflect our particular design (e.g., removing the 1"x5" slit requirement in favor of a requirement for areal coverage). Instead of a requirement for long-term stability, we require the ability to calibrate each spectrum to <3 m/s, with a <1 m/s goal. The spatial sampling on the sky is set by our spectral sampling rather than on an arbitrary plate scale. The field of view was increased to provide natural guide star availability and to provide future AO control of the deformable secondary.

4. INSTRUMENT CONCEPT

4.1 CONCEPT OVERVIEW

HROS will be a multi-purpose high-resolution optical spectrograph for the TMT. In order to meet its operational and functional requirements, we have several issues to confront. First, the large size of TMT requires either very large optics or small spatial sampling. Second, the telescope has to provide a spectrum from 310 nm to 1100 nm (ideally in one exposure to maintain efficiency). Second, the quality of the spectrum has to be high across the bandpass to allow science at arbitrary wavelengths. Ideally, the spectrograph will maintain high efficiency over as a large a bandpass as possible. Third, the instrument must maintain excellent imaging over the entire bandpass, which sets a strong requirement on the achromatism of components that see the entire waveband, such as the foreoptics, collimator, and focal plane components. Our solution is to use a multiplexed 1st order spectrograph, as illustrated in Figure 4-1. For this concept, components that have significant efficiency variations with wavelength (detectors and gratings) may be optimized at each wavelength to maximize performance. Broadband components can be either reflective (collimator) or very carefully chosen to minimize impacts (ADC, FIFU). The CU-HROS design is intended to provide high throughput, high stability, and high spectral resolution spectroscopy at optical wavelengths. This is accomplished by using an array of high performance dichroic mirrors to direct light into 32 narrow band first order spectrographs. The nominal design operates in a single mode with a minimum number of mechanisms, covering the wavelength range of 310 – 1100 nm at R=100,000 in a single integration.



During an observation, light from the telescope is directed to the Nasmyth platform and into the CU-HROS foreoptics. These include the ADC, derotator, stops, and baffles. A fast steering mirror is used to increase image stability, and the light then passes to the focal plane reimaging optics (if needed) to reform and possibly magnify the focal plane onto the fiber integral field unit (FIFU) array. Foreoptics also include a camera for acquisition, an absorption cell for wavelength calibration, and pickoff mirrors for the AO system.

After being directed into the fiber IFU (FIFU), the image is dissected in a coherent manner, and the 2x10mm entrance slit is reformatted to a 0.2x100mm pseudo-slit and. The FIFU includes additional fibers dedicated to wavelength calibration sources.

Figure 4-1: Schematic diagram of the CU-HROS Concept. Light from the TMT is fed into the CU-HROS instrument enclosure and reconstructed into a 0.1" pseudo-slit via a fiber optic IFU. The spectrum is divided into 32 spectral bins via a 31 element dichroic mirror array, feeding 32 narrow band, first order spectrographs with near constant resolution and efficiency.



Light from the FIFU is collimated and then passes through a series of dichroic mirrors in 5 banks, subdividing the spectrum into 32 channels of roughly equal spectral grasp. Each channel then feeds a single spectrograph bench operating in first order and optimized for its own narrow wavelength band. A possible physical layout is shown in Figure 4-2. This arrangement uses a single additional fold mirror to place the blue channels directly below the red channels to increase packing efficiency. Individual spectrograph benches could either be mounted vertically, or set up with the dispersion plane vertical, depending on the results of structural and thermal analysis.

By operating the spectrographs near Littrow in negative first order we eliminate positive and higher positive orders, substantially increasing the efficiency of the spectrographs. Careful design of the holographic gratings further optimizes throughput, and use of holographic gratings greatly reduces scattered light.

Figure 4-2: One possible layout of the dichroic array. Light enters from the upper left, and either passes through to the upper red arm, or is reflected back and folded into the blue arm (shown in green). The feed points for the individual spectrograph benches are indicated by spheres; the benches would either be vertical, mounted above and below the dichroic planes (requiring an additional fold mirror for each channel), or with the grating in the dichroic plane, and the dispersion plane vertical. The arrangement shown is not optimized, and is subject to more detailed analysis. Overall size is roughly 5x3x0.75 m without the spectrograph benches.

4.2 OPTICAL DESIGN

4.2.1 Design Overview

When we conceived of CU-HROS, we started with the conventional echelle-based design, which has long been the standard for high-resolution optical spectrographs for ground-based telescopes. The echelle has many advantages, in particular making use of the 2D format of the CCD for the high spectral bandpass at high resolution. However, as the telescope becomes large (such as 30 meters) and the instruments are seeing limited or



near seeing limited (as in the optical) the size of the instrument grows dramatically. For an instrument such as the one specified in the original HROS requirements — a slit-limited spectral resolution of 50,000 matched to a 1" slit — the beam size at the echelle is 0.9 m, requiring a very large echelle mosaic, with equally large optics at other points in the system. In addition, since an echelle works over a large bandpass, the variation in grating efficiency dominates the system throughput. The option to use a prism cross disperser becomes impractical due to the very large size of the diffracted beam, further reducing efficiency. Therefore, we decided to move in a new direction to solve these problems. Our design confronts the issues of very large optics by using many smaller optics and designing spectrographs around very narrow wavelength coverage while decoupling spectral performance from the telescope. In our current age of inexpensive silicon-based devices, it becomes cost-effective to consider the trade between large optics and multiple CCDs. To enable our design, we first decided to decouple the spectral performance from the seeing performance of the telescope. This required the addition of an integral field unit that uses 0.1 arcsec spatial elements, and restacking the fibers to form a pseudo-slit that is 0.1 arcsec wide. These changes allow us to retain all the light from the target while simultaneously allowing narrow spectral resolution elements. While superb images will improve our performance by requiring fewer pixels to be sampled in a given observation, decreased seeing will have no effect on our spectral resolution. To reduce our sensitivity to variation in grating efficiency over large wavelength regions, we replaced the single, multi-order grating of the echelle with multiple, individually optimized 1st order spectrographs. With dichroics becoming both less expensive and more efficient, we chose to split our light into discrete channels by the use of a dichroic tree to produce many channels through blue/red splits at each dichroic. This format allows each spectrograph to be optimized for a particular, narrow (~3%) bandpass. With such a narrow bandpass, anti-reflection (AR) coatings on detectors become highly efficient, and gratings can be optimized to work on blaze. Figure 4-3 shows a block diagram illustration of the optical design of our CU-HROS concept. Light from the telescope enters the foreoptics (ADC and derotator), and passes any pickoff mirrors or beam splitters required for guidance or AO wave front sensor (WFS) feed. The focal plane is reimaged onto the fiber optic integral field unit (FIFU) which reformats the 1 arc second point source field of regard into a long, narrow pseudo-slit: in terms of the 2mm/arc second plate scale, the pseudo-slit is approximately 0.1 arcsec wide by 10 arcsec long. The FIFU also allows for the injection of spectral calibration lines into dedicated fibers, allowing for continuous dispersion drift monitoring. Next, the light is collimated and directed into an array of 31 dichroic mirrors arranged in 5 banks which subdivide the instrument band pass into 32 relatively narrow bands. Each of these band enters an individual spectrograph bench with grating, camera



optics and detector optimized for the working wavelength range, allowing for uniformly high efficiency at all wavelengths.

Figure 4-3: Block diagram showing the major components of the CU-HROS design for a single

channel.

4.2.2 Major Optical Subsystems and Components

4.2.2.1 Foreoptics

4.2.2.1.1 Atmospheric Dispersion Compensator

All optical astronomical telescopes need to account for the chromatic aberrations induced by traveling through the Earth’s atmosphere at a non-zero zenith angle. In general, long observations will require continuously changing correction as the target moves across the



sky. To accomplish this, we will incorporate an atmospheric dispersion corrector (ADC) in CU-HROS. The ADC will consist of a series of low-dispersion prisms to correct for the chromatic dispersion introduced by the atmosphere. We base our design on the one presented in Schroeder (1987). For the baseline CU-HROS design, fused silica is required to provide high transmission at short wavelengths. For our baseline FOV (70"), we will require 500 mm prisms. While large, this size is available. Alternate designs of ADC units that diverge from simple flat prisms may significantly reduce the cost for this application and would be considered in a design study.

Figure 4-4 An example set of ADC prisms for the 70 arcsecond FOV that would be sufficient for our needs.

4.2.2.1.2 Derotator

As the TMT will be an Altitude-Azimuth telescope, a derotator is required for fixed instruments. The most cost-effective way to implement a derotator for CU-HROS will be a three mirror all-reflective design. The mirrors would be comparable in size to the ADC prisms (500 mm), giving similar cost/benefit regarding the CU-HROS field of view. The



three-mirror derotator will need to be slaved to the telescope pointing system. The derotator should be placed prior to the natural guide star pickoff to ensure compatible guide star motion.

Derotator Specifications Mirror Figure < 63 nm wavefront error per surface Mirror Surface Roughness < 2.5 nm rms surface roughness

Table 3: Requirements for the derotator

4.2.2.1.3 Beam Speed Optics

After the TMT telescope focus, with a field stop to bring the HROS FOV to the desired size (70 arcsec), we will reimage the target field onto the FIFU. The reimaging optics will need to change the beam speed to accommodate the FIFU entrance speed. By reimaging the telescope focus, an exit pupil will be formed, allowing baffling of the stray light from the image plane. We are considering several choices for the reimaging optics. Using a reflective reimaging system has the advantage of excellent chromatic performance, but will likely be slightly more expensive than an equivalent refractive system. Our preference would be to use an off-axis multiple-mirror (Korsch 1991) system to assure adequate performance. The optics would also provide an exit pupil to permit baffling or Lyot stops while maintaining the image quality provided by the TMT telescope. At this point, several different designs would meet the specifications. A detailed trade study in conjunction with vendors to determine the trades between different designs in terms of affordability and technical risk would be performed in the design study stage. The baseline design requires mirrors with diameters on the half-meter scale, with surfaces described by conic sections. These do not present significant manufacturing issues.

4.2.2.1.4 Pick-Off Mirror

We will have a positionable pickoff mirror at a true focus to have more light available for guiding (faint targets in field) or for AO (either natural guide star or laser guide star).

4.2.2.2 Target Acquisition and Guide Subsystems

4.2.2.2.1 Acquisition Procedure and Required Hardware



The baseline design specifies a fixed pattern of FIFUs with the option to upgrade to independently positionable multiplexing units. We will require a simple method for acquiring targets and guide stars. The acquisition process would be as follows: First, the telescope slews to the general field of the target(s) to within 5 arcsec. Second, a pickoff mirror is inserted at the foreoptics pupil, permitting the field to be observed by an acquisition camera for locations of the targets. Third, the movable fiber bundles will be positioned on their targets and the guiding system will be positioned on a guide star. Finally, the pickoff mirror will be removed and the observation can begin. We will require an acquisition camera and movable pickoff mirror at the foreoptics pupil. The acquisition camera will be a fast (f/1.0) final camera with modest imaging over a broadband coverage (to maximize ability to find faint targets). Since the light is sampled at a pupil, the full field will be available for the acquisition camera. The baseline guide system will be a system similar to a quadrant anode at the reimaged telescope image plane. The guider will be on a movable stage, as will the fiber bundles (similar to the Hectoechelle instrument on MMT). This allows arbitrary orientation of the guide star and scientific targets. If no suitable guide star exists in the field of view, we can leave the acquisition camera pickoff at the pupil and guide from the image of the target. The disadvantage is a reduced throughput in the science beam. While undesirable, this may be the only way to point at regions of very low stellar density. We have chosen the 70" field of view to ensure that a suitable guide star (of which many are available to a 30-m telescope) will be available in virtually every field.

4.2.2.2.2 Guiding Procedures and Required Functionality

Guiding procedures will be similar to those for existing instruments. The image of the guide star will be placed on the guiding camera. Drift offsets in centroid of the guide star will be relayed to the fast guide mirror (and ultimately, the deformable secondary mirror) or to the telescope itself to correct pointing during an observation.

4.2.2.3 Active and Adaptive Optics Subsystems

One of the principal advantages of the CU-HROS design is that the maximum spectral resolution/optics size is independent of AO performance. As a result, the principal improvement seen when AO is added will be in terms of efficiency and S/N as the seeing disk is reduced. However, the improvement in observing efficiency through the use of AO can be quite dramatic, particularly at red wavelengths. Models run by Dave Anderderson at HIA for GLAO implementation on HROS indicate an improvement in the FWHM of the half-light radius of the seeing disk of up to 0.15" in V and 0.3" in R. A



target of mAB=19 requires 9 hours of observing time to obtain a R=100,000, S/N=100 spectrum when the seeing disk is at 0.5"; the exposure time drops to 2.5 hours for 0.3". As a result, HROS should be designed to accommodate AO, if not at first light, then at some later point in its lifetime (when AO on TMT is sufficiently sophisticated to support optical AO, for example). There are multiple schemes to implement AO for HROS. The small field of view over which corrections are required ensures good performance from ground layer adaptive optics (GLAO), single laser GLAO (SLGLAO), or laser tomography schemes. We consider the SLGLAO scheme to be the baseline AO scheme for an incremental upgrade to CU-HROS. SLAGO provides adequate PSF improvement on axis, which is all that is required given the narrow field of view of HROS. However, it must be noted that any form of AO using a laser guide star may be problematic for CU-HROS, as the sodium guide star lasers operate at 589 nm, in the center of the CU-HROS band pass, vastly overwhelming any anticipated stellar signal in that region, and providing a potential (and nasty) source of stray light. In this respect, GLAO may be preferable to SLGLAO, because the lasers are off-axis in the former scheme. Use of the full constellation of GLAO lasers would require increasing the foreoptics diameter to accommodate an approximately 140" field of view, however, as compared to the baseline 10" field of view for HROS, and the 70" field chosen for CU-HROS to accommodate SLGLAO and NGS acquisition. Moreover, the large spot size of the laser guide stars (70" FWHM) ensures that even GLAO will feed an uncomfortably large amount of optical photons into the instrument (the lasers emit several billion photons per second per laser). We will provide the capacity for SLGLAO by reserving space for a single dichroic mirror to direct 589 nm light to a wavefront sensor via a variable optical delay and by designing the foreoptics to accommodate a 70 arc second field of view, as required for a single laser guide star. The design of the single wavefront sensor and delay will be dictated by the standard design selected for TMT. Additional dichroics may be required to sufficiently attenuate this light before it enters the spectrograph cavity. As light of this wavelength is in the middle of the HROS band pass, and all detectors will be sensitive to this light, extreme care must be taken to avoid allowing this light to enter the spectrograph cavity. Because the SLGLAO signal is so much brighter than any anticipated astronomical targets, even normally acceptable scattered light and off band rejection levels will allow sufficient out of band light into the various detectors to degrade observations. In order to accommodate the SLGLAO, we would need to insert a narrow band dichroic mirror to redirect the LGS light into a variable path length relay mirror system (trombone), and then to a wavefront sensor. This sensor would be the standard wave front sensor (WFS) selected by the project in order to leverage prior development, and to ensure compatibility with the AO system.



The cost of this system is listed as an optional upgrade. This is discussed further in sections 6.2.1 and 7.2.

4.2.2.3.1 Fast Guide Mirror (Tip-Tilt)

CU-HROS includes a tip/tilt mirror in the foreoptics to mitigate atmospheric and low frequency mechanical vibrations. This requires a fast sampling cadence and an natural guide star within the 70 arcsec field of view of the CU HROS. On average over the full sky, a 70" FOV contains at least one V=19 star (Cox 1999). For sightlines out of the Galactic plane, this FOV will contain at least one V=21 star, from which >10,000 photons/sec will be available in the optical waveband for guiding. Eventually, the TMT deformable secondary will have sufficient dynamic range to accommodate the tip/tilt correction required by CU-HROS. With the resolution element of CU-HROS at 0.1 arcsecond, the residual jitter at the 10 mas level is insignificant to our instrument.

4.2.2.4 Calibration Subsystems

The calibration system has a twofold purpose. First, the calibration subsystem must provide wavelength knowledge during a single observation and between observations separated by an arbitrary amount of time. Second, it must provide a flat fielding capability to improve S/N. Bias (zero) and dark frames will also be necessary for calibration and will be enabled by the system shutter(s). Flat fields are intended to provide a measure of the pixel to pixel variability of the detector sensitivity, and are generated by direct illumination of the CCD. This can be accomplished by flood illumination of the individual spectrometer benches, by feeding a semi-collimated source into each grating in such a location that 0-order light from the grating will be directed onto the detector, or by feeding continuum light into the FIFU entrance aperture. This has the advantage of mimicking the illumination geometry used by the science channel, and may provide a more reliable flat field. The flat field system must provide sufficient counts per resolution element to support S/N 300+ spectroscopy at the design resolution. We also recommend that TMT support the acquisition of daytime sky flats to provide sufficient photons at the bluest wavelengths covered. The CU-HROS instrument will include diffusing filters to support observations of the sky.

4.2.2.4.1 Wavelength Calibration

Wavelength calibration must be directly correlated to the observed spectrum, must provide adequate S/N during short flashes, and must provide a reasonable density of



calibration lines across the entire instrument band pass. We are considering three wavecal ‘sources’: The classic iodine absorption cell, a thorium/argon lamp, and a laser frequency comb. The first two are in use in existing telescopes, and the third appears to be readily implementable using current technology. Today, iodine absorption cells are standard wavelength references for high velocity precision observations (for planet searches), and have provided reliable and cross-system replicatable absorption spectra, but at the expense of depressing throughput of the target continuum by 50%. For bright targets, this sensitivity loss is offset by the advantages of introducing calibration features directly into the beam. An iodine absorption cell will be included in the CU-HROS concept to maintain continuity with current observing programs and to act as a calibration source for certain observations. CU-HROS will feature an iodine absorption cell that will either be deployable into the converging beam from the telescope before light enters the fiber IFU, or after the beam is collimated. Both locations have advantages and disadvantages: Placing the absorption cell in the converging beam will introduce chromatic aberration and require an offset to the location to the focal plane, but will not require as large of an absorption cell as placement in the collimated beam before the first dichroic. In addition to an absorption wavelength reference, CU-HROS will incorporate two additional wavelength sources: An emission line lamp (or set of lamps – some options include Th/Ar, He/Ne/Ar, and bright Ar, Ne and Xe ‘penrays’), and a filtered laser frequency comb. Both sources will inject narrow emission bands into dedicated fibers integrated into the fiber IFU so that emission spectra can be observed simultaneously with the science observation. Absorption spectra from the iodine cell will be used to verify that there is no systematic offset between the emission source locations and the science spectra. Th/Ar and He/Ne/Th lamps are dim but well characterized, and provide rich spectra in the region of interest. For example, Th/Ar lamps provide over 4000 stable, identified lines between 300 and 1100 nm, and have provided excellent reference spectra for many instruments. For the highest resolution, highest velocity precision observations, however, these lamps are beginning to dominate errors on some instruments (such as the planet-finding instrument HARPS): arc lamps suffer from line blending, somewhat poor identification of features, and dramatic fluctuations in line strengths over the full waveband. As a result, we have also looked into other options for wavelength calibration sources. In addition to emission line spectra, recent work in at the National Institute for Standard Technology (NIST) in Boulder, Colorado, has produced laser frequency combs with the



potential for use as extremely stable, well characterized line sources. A frequency comb produces emission lines (‘teeth’) at a regular frequency spacing: ωn = nωγ + ωo, where ωγ is the driving frequency, and the offset frequency ωo can be phase locked to a highly stable atomic source (see Figure 4-5).

Figure 4-5: Consecutive pulses emitted by a mode-locked laser (a) and the corresponding

spectrum(b). (from Udem, et al., 2002)

Current work has generated line spacing as wide as 2GHz (approximately 250,000 resolution). This line spacing can be increased by using an etalon to filter the unwanted lines, reducing the number of lines by an order of magnitude. The advantage of this light source is that it produces regularly spaced, extremely stable emission lines across the instrument band pass: the driving frequency and offset frequency can be phase locked to an off the shelf atomic frequency source. This results in wavelength accuracies as great as 1 part in 1010 when coupled to a rubidium microwave transition source. Even though the output intensity varies over several orders of magnitude, the output is very intense relative to our requirements (outputs are on the order of µW/tooth, while we require fW), so selective filtering could be used to normalize the beam. Although this technology has yet to be employed as an astronomical wavelength standard, there are no foreseeable technological roadblocks to implementation. High repetition rate frequency combs are readily fabricated and economically viable for use on a project of the scale of the TMT. Output from the frequency comb could be shared with multiple instruments and the unfrequency-filtered signal would make an excellent pseudo-continuum for flat-fielding.



4.2.2.4.2 Photometric Calibration

Photometric calibration will not be internal to the instrument, but will be provided by the observation of standard stars. Because HROS will provide simultaneous spectroscopy over the full optical waveband, it will be a powerful spectrophotometer for a number of science cases, such as stellar modeling analyses and observations of active binary systems.

4.2.2.5 Focal Plane

While the default focal plane for HROS is only 10x10 arcsec, with a 5x1 arcsec entrance slit, we will equip the CU-HROS with foreoptics capable of supporting 70 arcsec field of view in order to allow for acquisition of guide stars and to accommodate SLGLAO. Other features of the focal plane include space for dichroic mirrors to redirect and attenuate the SLGAO beam, and a deployable pickoff mirror to support acquisition, guiding and tip/tilt correction. The most critical element of the focal plane is the fiber integral field unit, which is used to remap the entrance aperture into a 0.1 arcsec pseudo-slit.

4.2.2.5.1 Fiber Integral Field Unit

CU-HROS relies on reformatting the 1 arcsec entrance slit to an effective 0.1 x 10 arcsec pseudo-slit. The current design relies on a fiber optic IFU (FIFU) to perform this image dissection. The design could be achieved with a conventional image slicer, but at the cost of reduced performance; the trade study comparing these options is discussed in Section 6.3. Performance of fiber IFUs have been limited in the past by the fact that these devices have been retrofitted onto existing spectrometers. As a result, they have been subject to significant vignetting losses within the existing instrument, and have suffered reduced resolution, presumably because the inherent focal ratio degradation (FRD) of the fibers results in a much faster beam than the existing spectrograph optics were designed to accommodate. Since our design will be based from the start on an extended, fast source, we do not anticipate substantial efficiency loss from vignetting, or resolution loss from the fast output beam. The initial motivation for introducing an integral field unit was our desire to reduce the width of the entrance slit to no more than 0.10 arc seconds equivalent plate scale in order to reduce the size of our optics. Given an f/15 entrance beam, this factor of 30 reduction in the image size is not physically possible, as it would violate the LaGrange invariant. While only a modest reduction in magnification is needed to avoid this violation final imaging optics would be too fast to be practical.



We can avoid this conflict by restructuring the beam with an image dissecting integral field unit. Rephrasing the problem in terms of the 0.5 arc second per mm plate scale, we would take the 2x2mm entrance aperture and demagnify and remap it to roughly 13.7 by 0.066 mm. We accomplish this by placing an array of 91 0.066 mm core fibers at the focal plane on 0.2mm centers, and remapping that to a vertical array one fiber wide. For a 2x2 mm rectangular array this would result in 100 66 micron spots in a vertical array, separated by roughly 150 microns center to center (including cladding and mounting hardware). The first problem with this approach is that the fill factor for these fibers is quite poor, so some sort of microlens or microcondenser array is needed to feed the individual fibers. This leads to the second difficulty: a rectangular array introduces significant geometrical focal ratio degradation, and should be avoided if possible. This problem is resolved by using a hexagonal packing geometry and tessellating the microlens array. Excellent discussions of microlens coupled fiber IFUs can be found in Kenworthy (1998), Kenworthy, et al. (2001), and in Allington-Smith, et al., (2002). In Allington-Smith, et al. (2002), they describe an IFU for the Gemini Multiobject Spectrograph similar to the instrument we propose, and provide performance data based on actual operations: 68%, including losses from the spectrograph stops. They suggest that their design could achieve efficiencies as high as 79% if there were no constraint on the angular extent of the output beam. Since we are designing our spectrograph collimator around the fiber IFU design, we do not expect to be constrained by light loss at internal stops. Our anticipated efficiency is based on previous instruments, assuming 74% efficiency. This is at the high end of fiber IFU efficiencies, but is not unreasonable since it does not include any internal baffle losses. As part of our end to end optical design we are carrying vignetting in the system as a separate line item. While the default design for HROS called for a rectangular entrance aperture for a point source spectrograph on a telescope with circular diffraction symmetries, we will build an array of several 1 square arcsec hexagonal fiber bundles. In order to take advantage of improved performance when the seeing disk decreases, we are considering using a methodology similar to that described in Kenworthy et al. (2001), where the microlens to pseudo-slit mapping spirals outwards from the center of the microlens array (Figure 4-6). This would concentrate light towards the center of the output array, would reduce the number of illuminated pixels at the detector, and would allow on-chip binning at the detector to reduce detector read noise. This and other mapping schemes will be evaluated prior to implementation.



Figure 4-6: (a) Illustration of one possible remapping of the fiber optic IFU input to the pseudo-slit output. Tessellated hexagonal microlenses (used to minimize geometrical FRD) are used to concentrate light onto the individual fibers, which are then remmaped into a long, narrow pseudo-slit. Various remapping schemes, such as the spiral mapping shown here, and various raster mappings will be considered prior to implementation. (b) Distribution of light at the pseudo-slit for a 0.5 arc second Gaussian PSF. Spiral image dissection concentrates light towards the center of the pseudo-slit.

Rather than a fixed contiguous 1x5 arcsec slit, we would provide an array of 5 to 7 1 square arcsec IFUs in a nominally fixed grid. This would allow for sampling multiple points on an extended object, and for a well separated sky channel. Alternatively, the central IFU could be larger, perhaps 2 or 3 square arc seconds, with fewer sky channels in a fixed array around that. This design allows for a straightforward incremental upgrade path to a movable sensor array by placing the non-central input arrays on 2 dimensional translation stages. There are additional advantages to using a fiber IFU for the input. One is the ability to translate the output to a fraction of a resolution element in order to perform focal plane splits if it is determined that this would increase the instrument S/N. Another advantage is that we can insert dedicated calibration fibers into the IFU and inject the calibration signal directly into the spectrograph while science spectra are being recorded. For example if two fibers at the top and bottom of the pseudo-slit were dedicated to calibration, then the signal from these would be spatially separated from the science signal, but could be correlated to the science signal with very high repeatability since they are being emitted from the same physical structure. Finally, the output geometry of the FIFU can be tailored to the needs of the collimating optics

4.2.2.6 Collimating Optics



We will use an all-reflective collimator based on a three-mirror anastigmat (Korsch 1991) to provide an achromatic collimated beam without residual astigmatism or coma. A three-mirror anastigmat provides control over coma, spherical aberration, and astigmatism with an arbitrary field curvature. In addition, the collimator cannot leave aberration in the beam to preserve the camera performance post-grating. With the freedom to place the fiber ends at arbitrary positions, a curved input plane is easily tolerated, reducing the constraints of the collimator. Further optimization is incomplete regarding the pupil location as it applies to the size of the dichroic tree. We expect the optics to be ~1 meter class. With advances in mirror polishing technology this does not represent a significant increase in cost, and for larger optics would be a cost savings relative to transmitting optics. In addition, the use of reflective surfaces instead of transmission will minimize chromatic aberrations throughout the system. In Figure 4-7 we present a preliminary concept for a collimator suitable for CU-HROS. The design uses off-axis elements to avoid vignetting the light from the fiber bundle. At present we are also considering other designs (4 mirror examples) that produce a final pupil further from the collimator to keep vignetting effects to a minimum.

Figure 4-7: Example collimator based on the formalism of Korsch (1991). The collimated beam exits

at the lower left side of the panel.



4.2.2.7 Dichroics

Dichroic mirrors provide the primary wavelength selection for the CU-HROS instrument, and the performance of the instrument is dependent on these mirrors operating at high efficiency. Dichroic mirror technology has undergone substantial advances in the last several years, yielding mirrors with sharp transitions (>5nm from 90% rejection to 90% acceptance) and high efficiency (typically >95% transmission/reflectance). Measured performance of dichroic mirrors already produced by Barr Associates support this contention (Figure 4-8 and Figure 4-9). The ability of Barr Associates to deliver optics with the performance shown in these two figures supports our confidence in their ability to meet our performance requirements. Note that the fluctuations in transmission/reflection shown in these curves also appears (to a lesser extent than for the WFC3 data in Figure 4-8) in the theoretical data for the HROS dichroic mirrors; however, the locations of the dips in performance vary from filter to filter, and tend to average out after 5 filters (see Figure 4-11).



Figure 4-8: Measured performance for Wide Field 3 dichroic mirror F606W (blue line) delivered to NASA. An earlier, rejected mirror is shown for comparison (dashed red). Figure from WF3 monthly status review presentation, Sept. 1 2005 (John MacKently presenting).

Figure 4-9: Measured performance for an existing 125x125 mm dichroic mirror; data provided by Barr Associates.

Figure 4-10 shows a simulated spectrum of a QSO before and after passing through the dichroic tree. Barr expects to exceed our specification but for the purposes of this illustration we have reduced best performance to our requested specification, no more than 95% transmission/reflection/out of band rejection. Different colored bands represent the light entering each of the 32 spectrograph benches. Note that the bright feature near 800 nm spills into the adjacent channel in this worst case model. Low scatter holographic gratings (typically better than 10-5/Å )will allow this out of band light to be efficiently rejected.



Figure 4-10: Simulation of a typical AGN spectrum before and after passing through the dichroic tree, based on Barr theoretical data for our requested mirror set.



4.2.2.7.1 Component Description

Each optic is a single-transition dichroic, with transition wavelengths selected to allow for a binary wavelength sorting into 32 relatively narrow bands. A sample array and the corresponding channel wavelength ranges are given in Table 4. Performance of the individual channels is shown in Figure 4-11 and is based on theoretical data provided by Barr Associates, with performance degraded to a maximum reflection/transmission of 95%, and a maximum out of band rejection of 95%. Note that these wavelength ranges are only provided for illustration and are based on model data provided by Barr associates for the initial concept study. Final transition wavelengths would be carefully selected to avoid placing known, interesting astronomical features in any of the interchannel overlaps. The dichroic mirrors will have a 200 mm clear aperture, and be optimized for 15° or less angle of incidence in order to minimize polarization effects. Ideally, each mirror would have less than λ/10 wave front error for both the reflected and transmitted beams, but this may prove impractical in the reflection case, so wavefront correction optics may be necessary for some channels.

Sample HROS Dichroic Edges and Corresponding Channel Widths

Dichroic Transitions HROS Channel Limits

Dichroic ID

Transition Wavelength

(nm) Channel Wavelength Range

(nm) 0-1 574.5 0 300.0 313.0 1-1 413.5 1 311.9 326.0 1-2 792.5 2 324.9 339.0 2-1 352.0 3 337.9 352.5 2-2 485.5 4 351.4 367.0 2-3 673.0 5 365.9 382.0 2-4 352.0 6 380.4 397.5 3-1 325.0 7 395.9 414.5 3-2 381.0 8 412.4 431.0 3-3 447.5 9 429.4 448.5 3-4 526.5 10 447.4 467.0 3-5 619.5 11 465.4 486.5 3-6 730.0 12 484.4 506.5 3-7 861.5 13 504.4 527.0 3-8 1017.0

14 525.9 549.5



4-1 312.5 15 547.4 573.5 4-2 338.0 16 570.4 596.0 4-3 366.0 17 593.9 620.5 4-4 396.5 18 619.4 647.0 4-5 430.0 19 644.4 674.5 4-6 466.0 20 671.4 702.0 4-7 505.5 21 699.9 731.0 4-8 548.0 22 729.4 762.0 4-9 594.5 23 759.9 794.5

4-10 645.5 24 791.4 827.5 4-11 700.5 25 825.4 862.5 4-12 761.0 26 860.4 898.5 4-13 826.5 27 896.9 938.0 4-14 897.5 28 933.4 976.5 4-15 975.5 29 974.9 1018.0 4-16 1060.0 30 1015.9 1061.5

31 1058.9 1120.5 Table 4: Sample HROS Dichroic Edges and Corresponding Channel Widths

Figure 4-11: Performance of filter stacks for each channel based on theoretical data provided by Barr Associates. As in the previous figure, predicted performance is reduced to match specification. Roll off at short wavelengths is induced principally by the short wavelength performance of the initial filter,

which must have excellent performance over an extremely wide wavelength range. Also note that high frequency fluctuations in individual filters do not correlate, but instead tend to average out after 5

transmissions/reflections.



4.2.2.7.2 Tolerancing

In the case of dichroic vibration, we can tolerate 0.1 arc seconds rms vibrational motion about one axis, and are, for all practical purposes, unconstrained in the other (based on a preliminary tolerance distribution). The constrained axis depends on the orientation of the diffraction gratings. For the baseline design with the dispersion plane perpendicular to the dichroic fold plane, and the dichroic mirrors arrayed in the horizontal plane, the tightly constrained rotation is rotation about a horizontal axis. Root sum square (RSS) combinations of these motions would contribute blur over 0.2 resolution elements. If this motion is vibrational, then the it would only contribute to a line broadening. If thermal, it could introduce a line motion, which would have a much more severe impact on science return. Thermal motion introduces line motion principally for fiber IFU displacement in the dispersion direction, and for transverse motions of the detector relative to the diffraction grating. Clever design will allow these elements to be metered to common structures with little offset in the critical dimension. However, of greater concern are the second order effects of rotations induced by thermal gradients in the dichroic mirror mounts. We will want to stabilize the optical cavity to less than ±1°C, and thermal gradients must be kept to a minimum. The extent of thermal rotation is difficult to estimate at this level of development, but a crude model suggests that thermal gradients across optical mounts of order 0.1°C will affect instrument performance. This constraint will be defined during the instrument design definition phase.

4.2.2.8 Dispersing Optics

Holographic gratings are the preferred option for CU-HROS. Holographic gratings are produced by the interference pattern of two stigmatic (laser) source point sources on a photo-sensitive material. The photo-sensitive material is developed and a series of smooth grooves are produced on the surface of the substrate. Due to the smooth nature of the grooves, holographic gratings have very low scatter (measured at better than 10-5 – 10-6 Å for the Cosmic Origins Spectrograph gratings). Our preferred option is to request a pseudo-sinusoidal groove profile by adjusting the exposure time to produce deeper grooves. This produces high efficiency devices (~60%), although over a relatively narrow bandpass (~5%). With the design for CU-HROS, each channel represents a 3% bandpass and very high groove efficiency can be achieved at all wavelengths. Suitable mechanically ruled gratings can be produced, but have significant disadvantages. They will always have much higher scatter due to surface roughness induced by the ruling engine. In particular this will affect channels (λ > 650 nm) with bright OH lines in conjunction with the spectra of faint targets. The high line densities we require in the blue channels may be difficult to rule for existing ruling engines (see the line densities in



Table 6). In our experience, high-quality mechanical gratings tend to be more expensive than equivalent holographic gratings, particularly at the high line density that CU-HROS would require. When 32 unique gratings are required, the cost-effectiveness of the holographic gratings represents an even larger advantage. The mechanical gratings’ efficiency will be comparable to holographic gratings, but for the level of robustness and precision required, holographic gratings are the preferred option. We have been in discussion with Jobin-Yvon, a world leader in holographic grating technology. They are very interested in creating such large-format gratings and in helping us to develop the appropriate groove profiles to maximize efficiency.

4.2.2.9 Camera Optics

The cameras required for CU-HROS are fast (f/1.8) cameras with wide fields (14 degrees) and excellent aberration control (spot FWHM 15 microns). Preliminary investigations imply that several extant solutions meet our needs. In discussions with Zygo Corp and Goodrich Space Optical Branch, both vendors felt that appropriate, cost-effective solutions exist and could be produced. A significant advantage to the CU-HROS spectrograph cameras is the narrow wavelength range required for each camera, simplifying design by removing the need for achromatism. Our current baseline is to use a modified double-gauss camera that produces adequate (~20 micron) images across a flat field. Our camera design is aided by the significant anamorphism in the beam from the grating to the camera. By compressing the beam in the spectral direction, the final beam is closer to f/5, increasing performance. The trades that need to be examined closely in the future would be the relative advantages of one all-reflective system (simplifying overall system engineering) versus the slightly more complicated transmissive system that is our current baseline. The transmissive system does not have the broadband performance to be used in all channels; however, perhaps as few as three different configurations will be needed to accommodate the broadband performance of CU-HROS. Figure 4-12 is an example camera suitable for our design. The camera is based on a double-gauss system that produces <15 micron images without considering tolerancing or manufacturing error. The camera itself uses two aspheric elements, which will add to the cost of manufacture, although recent advances in figuring technology will mitigate this cost. The optics are in the 200 mm size range, increasing availability and reducing cost and risk to the program. We use the anamorphism present from the grating to help control aberrations in the spectral direction to smaller than the spatial on-chip binning size (5 pixels).



Figure 4-12: Example of a six-element double-gauss camera suitable for CU-HROS. Channel 17 (569 - 594 nm) uses two aspheric elements (first and last lenses). Note anamorphism in spectral

direction. Lenses are on the order of 250 mm.

4.2.2.10 Detectors

The detectors for CU-HROS represent conventional technology in a large format application. Commercial CCD vendors (e2v, Fairchild) have expressed interest in working with us to produce a series of custom CCDs for our application. Our current design would use 12.5 micron pixel, 9K by 3K CCDs, which will fit on a conventional six-inch silicon wafer. The well depth of such CCDs would be on the order of 100,000 electrons for the 12.5 microns pixel pitch. While there is a trade between larger pixels and reduced read noise, the resulting cosmic ray rate may be too high. At the faint limit for a seeing limiting telescope read noise dominates all noise sources. As we discussed previously, this is a function of the aperture and will affect all instruments equally: assuming a best-case F/1.0 final beam speed, the spatial scale of a 30-meter telescope is roughly 150 microns per arcsec, which for 15 µm pixels translates to >10 pixels per resolution element independent of instrument design.



In order for the reader to understand the number of pixels that CU-HROS would use compared to a similar conventional echelle spectrograph (with 15 micron pixels) we show illuminated pixels for different observing conditions and spectral resolutions in Table 5. Note that we used the default HROS concept from the SRD for the comparison (i.e., a 1 arcsec slit corresponds to a resolution of 50,000 on the detector). A resolution of 100,000 is achieved in the slit-based design either by using a 0.5 arcsec slit or by using an image slicer, depending on observing conditions. We also include the ability to profile the source over the slit to reduce the number of pixels sampled for the conventional echelle. CU-HROS by comparison is baselined for 100,000 spectral resolution for all seeing conditions. Note that because the CU-HROS concept retains all focal plane information, the choice of how many pixels to sum for the output spectrum is user-selectable (i.e., under poor seeing conditions, the user can choose to reduce read noise by extracting only the central 38 pixels at the expense of reduced target photons).

Observing Condition CU-HROS Conventional Echelle With Profile 1 arcsec, R = 100,000 156 200 160 0.5 arcsec, R =100,000 38 100 80 1 arcsec, R = 50,000 312 100 80 0.5 arcsec, R = 50,000 76 100 80 1 arcsec, R = 20,000 780 250 200 0.5 arcsec, R = 20,000 190 250 200

Table 5: Pixel Illumination for CU-HROS and a Conventional Echelle on TMT

The significant reduction in pixels used by CU-HROS at high resolution is due to the sampling of the 2-D point spread function on the fibers, rather than reducing the number of pixels along the slit. Note that the conventional echelle does illuminate fewer pixels for lower resolution modes (not including on-chip binning), but in the high resolution mode CU-HROS has fewer illuminated pixels, particularly in the case of good seeing where the CU-HROS advantage is clear. At worse than one arcsecond seeing, the conventional spectrograph would lose a significant amount of light due to slit losses when compared to CU-HROS with an oversized IFU (see 4.2.2.5.1) that can accommodate the broader PSF, albeit at the cost of more illuminated pixels. How do we address the increased number of pixels and consequent read noise increase over a conventional echelle for low spectral resolution modes? We have two possible solutions and will continue to look into others. The simplest would employ on-chip binning in the spatial (and spectral) directions to retain full resolution. CU-HROS provides a linear spectrum with only minor optical



distortions resulting in the ability to bin a significant number of pixels along the spatial direction before compromising spectral resolution. In particular, if lower resolution spectra are desired, binning in the spectral direction further reduces the effect of read noise. The user would be free to determine when on-chip binning is appropriate and the extent of the binning, based on information provided by the instrument team. In our conversations with vendors there are very few difficulties in binning up to 53 pixels together. Our pixels are large enough that the well depth is good even for a large number of pixels, particularly as the background is spread over nine times as many pixels as a ten meter telescope. The technique of on-chip binning will increase the amount of incorporated background automatically. The amount of binning on chip will be determined by the user and his or her acceptance of lost spectral resolution information. For a cosmic ray rate similar to Mauna Kea, 2.1 CR cm-2 min-1 (Smith et al. 2002) the mean time between cosmic ray hits on a given pixel is 320,000 minutes. If one can tolerate a loss of 1% of the total data due to cosmic rays, the maximum integration time is several hours without on-chip binning. For observations lasting one hour, the number of pixels that can be binned together is 53. Since there will be multiple exposures on a given target, the 1% loss per exposure will be insignificant as the specific data lost in one observation will be present in others. We would expect the observer will bin on-chip in the spatial direction for most observations and in the spectral direction only when running the system at reduced resolution. We estimate that 5 pixels in the vertical direction would be a typical binning scheme without compromising the sampling of the spot. For reduced resolution, we can bin on-chip in both directions, and the well depths (in particular for faint targets) should not present a significant issue. Note, though, that the details of the well depths in the read register will not be known until further into the design process because the specification for the CCD fabrication will be developed in conjunction with the vendor. The user would be allowed to choose the maximum integration time, although there will be recommendations from the CU-HROS team. A potential alternative technology that we have identified is the use of photon-counting CCDs. e2v has a new type of CCD becoming available that uses a series of avalanche diodes to multiply electrons on the chip prior to read. With a faint target, this will result in a noiseless read at high speed. Bright targets (photon rate close to the pixel read rate) will require modification of the read scheme. The modifications will increase read noise or slow read time; this may not present problems given that bright targets are dominated by photon noise rather than read noise. This technology is promising for astrophysical applications, and should be considered carefully in the design study phase.

4.2.3 Performance



4.2.3.1 Resolution

Each spectrograph is designed to provide similar resolution, with the bandpass for each spectrograph bench increasing with wavelength. Our baseline design is for 32 structurally identical spectrograph benches. Each spectrograph channel will have the same bandpass in terms of individual resolution elements so that the resolution will be close to constant across the entire wavelength range (ranging from R = 100,000 to 104,000). The following is a table of the wavelength coverage by channel. Also note the overlap between channels, preventing loss of spectral coverage and increasing overall efficiency.

Channel λstart

(Å) λend (Å)

Grating line density (grooves/mm)

1 3000 3130 6023 2 3120 3255 5791 3 3245 3386 5568 4 3376 3522 5352 5 3512 3664 5145 6 3654 3812 4945 7 3802 3967 4752 8 3957 4129 4566 9 4119 4297 4387 10 4287 4473 4215 11 4463 4656 4049 12 4646 4847 3889 13 4837 5047 3735 14 5037 5255 3587 15 5245 5473 3445 16 5463 5699 3308 17 5689 5936 3176 18 5926 6182 3049 19 6172 6440 2927 20 6430 6708 2810 21 6698 6989 2697 22 6979 7281 2589 23 7271 7586 2485 24 7576 7904 2385 25 7894 8236 2289 26 8226 8583 2196 27 8573 8944 2108 28 8934 9321 2022



29 9311 9715 1941 30 9705 10125 1862 31 10115 10553 1786 32 10543 11000 1714

Table 6: Spectral coverage for CU-HROS over the optical waveband

4.2.3.2 Efficiency

We have estimated the throughput of the CU-HROS instrument based on data provided by vendors or common industry capabilities, and the resulting throughput calculations are summarized in Table 7 below. Variation in efficiency is expected to be dominated by variations in the dichroic array throughput. Figure 4.7 shows the expected throughput modulation due to variations in the dichroic mirror performance. For components before the dichroic tree, we report the minimum anticipated performance. For components on the individual spectrograph benches, such as the camera optics and gratings, we anticipate uniformly high efficiency since each component can be optimized for the narrow operating band of each spectrograph bench. The fiber IFU efficiency is based on instrument performances reported by Allington-Smith, et al. after removing the effects of vignetting introduced when these IFUs were retrofitted onto existing instruments, and reducing the performance by a safety factor of 5%. Since we are designing our optical path around the characteristics of our IFU output (and tailoring the IFU itself to accommodate the needs of the optical system) we do not anticipate any impact on the performance of the optical system due to the IFU pseudo-slit length or to the inherent focal ratio degredation of a fiber-optic system.

Top Level Efficiency by Subsystem System Efficiency Fore optics 0.767 FIFU 0.740 Collimator and vignetting 0.867 Dichroic tree 0.774 Spectrograph Bench 0.498 Net HROS Performance: 0.190

Table 7: Estimated throughputs for CU-HROS

4.2.3.3 Stability



Spectral stability of the HROS instrument can be broken down into two regimes: long term and short term, with long term referring to stability between observations (CCD integrations), with two observations separated anywhere from hours or years, and short term referring to any motions occurring more rapidly than a single integration. Sources of short-termm spectral motion could include component and inter-component thermal drift, optic jitter (especially in the case of the dichroic mirrors), and atmospheric diffraction within the instrument due to turbulence and convection. We have modeled inter- component motions and jitter, but have not explored turbulence and convection because the baseline instrument design calls for all optics after the fiber IFU to be placed in vacuum to minimize acoustically induced jitter in the dichroic mirrors and to increase the thermal stability of the instrument. Long term spectral motion will be monitored, rather than eliminated, through the use of real time calibration systems. This is discussed in section 3.1.6. Photometric stability is also desirable, but not as critical. Long term degradation must be avoided by limiting exposure to ambient atmosphere and maintaining strict clean room procedures for the optical cavity. This is facilitated by having the optical cavity open directly into a permanent clean room mounted on the Nasmyth platform with CU-HROS.

4.3 MECHANICAL DESIGN

The baseline design for the CU-HROS is a largely static instrument with virtually no active mechanisms after the foreoptics. The dichroic mirrors will be mounted on a common optical bench in two banks, with the lower bank folded under the upper in order to increase packaging efficiency. Individual spectrograph benches sit above and below the planes of the dichroics to further minimize throw length and overall volume. CU-HROS must provide a high degree of spectral stability, wavelength knowledge, and repeatability. This degree of stability will require placing the instrument in a thermally and acoustically isolated enclosure. The present design baselines a vibration-isolated bench in a thermally controlled vacuum chamber. This effectively isolates the instrument both from ambient thermal fluctuations and from acoustical vibrations. A vacuum chamber would provide an excellent degree of thermal and acoustic isolation, but we would perform a trade study in the design study phase to determine if an adequate degree of isolation can be achieved without the added mass and expense of a vacuum chamber. This is discussed in greater detail in Section 8.6.

4.3.1 Physical Location and Dimensions



Aggressive packaging of the CU-HROS instrument (including folding the blue channels under the red) as shown in Figure 4-2 would compress the footprint to at most 5x3x3 meters for optics (this includes folding light 90° for the spectrograph benches). Foreoptics could be accommodated in the same enclosure, or could be isolated for maintenance. Overall footprint of the instrument enclosure will be no more than 7x5x4 meters. With the outer enclosure/clean room, the footprint will grow to no more than 9x10x4.5 meters. Optimization of the instrument layout may reduce this even further. Mass for CU-HROS will be approximately 45,000kg, with roughly half of that attributable to the vacuum chamber.

4.3.2 Enclosure

The outer enclosure for CU-HROS will consist of a 9x10x4.5 meter clean room designed to minimize acoustic coupling from the outside to the inner optics cavity. Inside of this room is the instrument enclosure which provides further acoustic, vibration and thermal isolation of the instrument.

4.3.3 Component Parts

CU-HROS can be broken down into several major assemblies:

• Primary enclosure (clean room) • Secondary enclosure (vacuum chamber) • Foreoptics housing (possibly a secondary vacuum cavity) • Optical systems • Foreoptics bench • Focal plane assembly (including FIFU) • Primary optical bench and dichroic Tree • Individual spectrograph benches (32)

4.3.4 Power Requirements and Output

Power requirements for the CU-HROS, based on values from existing components, are summarized in Table 8.

HROS Power Requirements Component Power Clean room 18.4 KW Vac Chamber 30.6 KW Vac shroud 10.0 KW Detector cooling 16.0 KW



Total 75.0 KW Table 8: Power Requirements for CU-HROS

4.4 ELECTRONICS DESIGN

The electronics design and components will assure that the instrument can be monitored, operated and commanded easily and reliably and will meet all of the requirements as defined in the HROS IFPRD as well as any observatory conformance requirements for the interface, power consumption and dissipation in the dome and safety operations.

4.4.1 Mechanisms

The baseline design has most of the components of the instrument residing inside a chamber. Each mechanism will be fabricated in accordance with vacuum operations for cleanliness, heat dissipation and lubrication. If it is determined that a vacuum chamber is not required, we will revisit these requirements in order to reduce cost and fabrication time (see Section 8.6). The list below shows the various systems and mechanisms currently identified. Atmospheric Dispersion Corrector (ADC): will consist of a rotating prism and will require a motor, an encoder and limit switches to move and monitor its position. This system will require the pointing information from the Telescope Control to properly align the prism. Image Derotator: will consist of three flat mirrors mounted in a rotating assembly. This system will require a motor, an encoder and limit switches to move and monitor its position. This system will require the pointing information from the Telescope Control to properly rotate the assembly as the telescope tracks the sky. Fast Steering Mirror: will consist of a fast (≥1 KHz) tip-tilt mirror placed in the optical path to remove the low order effects of seeing and any vibrations of the telescope/instrument. This system will use a natural guide star as the source and will have either a dichroic or pickoff mirror feeding a camera and a small embedded system to quickly acquire images, measure the star centroid and feed the required corrections to the tip-tilt mirror. An assembly with a series of neutral density filters will also be implemented to give access to a larger number of natural stars. Auto Guider: will consist of a pickup mirror that can be positioned in the field of view, which will feed the guide camera. The displacement of the guide star in the camera will be fed back to the Telescope Control System according to the format defined by the TMT



TCS team. The system will consist of an X-Y stage to position the pickup mirror with the associated camera and electronics. The guider will contain a set a neutral density and/or color filters to accommodate the wide range of potential guide star brightnesses. A frame grabber video card will reside in the Instrument Control Computer (ICC) to acquire the image, determine the centroid and send the information back to the Telescope Control System (TCS). Adaptive Optics (AO): will consist of a wavefront sensor (WFS) and an embedded system to analyze the images from the AO and pass the information to the TCS in the expected format so that corrections can be applied to the deformable mirror. The details of the requirements of the instrument’s AO capabilities will be determined at a later time to meet the requirements the TMT AO team will define. Shutter: will consist of either a single shutter at the entrance aperture of the instrument or a series of 32 synchronized high speed shutters in front of the entrance window of each CCD. The decision will depend on stray light analysis and trade studies of each implementation. The shutter(s) will be controlled by a single computer auxiliary electronics board. Grating Mechanism: will consist, for the baseline design, of a rigid mount with no allowable motion. The different instrument design options (low and very high resolution) will require either multiple gratings on the same mount or the ability to rotate the grating to scan the full wavelength coverage. These rigid mounts would then be replaced by vacuum rated rotational stages or rotational mechanisms with a fixed number of positions that can be locked in place. Fiber Positioner: will consist, in the baseline design of 5 fixed fiber bundles with no moving parts. The upgrade design includes a X-Y positioning mechanism for 4 of the 5 bundles to be positioned around the central fiber bundle. Calibration Platform: will consist of flat field and spectral lamps with a lamp select mirror, a lamp injection mirror, and fiber feeds. Those mirrors will be mounted on simple flip mechanisms and be operated with simple vacuum rated DC or stepper motors and limits switches. Iodine Cell: will consist of of a linear stage to move the cell in and out of the optical path. Standard vacuum rated stepper motors and a linear stage with relevant limit switches will be part of this system. The WFS and the auto-guider will relay their information back to the telescope at the rate and in the format expected by the TCS. Specific requirements have yet to be defined.



The commanding of the moving mechanisms will be done through TCP/IP connection to the various devices from the ICC. The ICC will also communicate with the various embedded systems (AO, auto-guider and the fast steering mirror) to initialize, start or stop the specific system. The communication from the ICC to the TCS will be according to the standards yet to be defined by the TMT TCS team.

Figure 4-13: Diagram of the different mechanisms and their links

Because of the clear separation of tasks between the wavefront sensor, the fast steering mirror, and the rest of the instrument’s functions (including the autoguider), it is natural to separate and assign these tasks to different dedicated embedded systems. The wavefront sensor, because of its fast image acquisition, intensive computation, and fast feedback to the TCS, will require its own embedded system. The same goes for the fast steering mirror but this will be a closed system with no feedback to the external world. The rest of the functions will be handled by the instrument control computer (ICC). The auto guider requires a much lower processing rate and will be handled by the same general purpose computer.



All data files will be saved in a FITS format with keyword-value pairs corresponding to all of the housekeeping information regarding the HROS instrument, the telescope, the current weather, and any other auxiliary instrument (guiding, tip-tilt, AO, etc.). If the power dissipation in the dome from HROS is above an acceptable level, the system will be vented to other parts of the building to minimize dome seeing effects.

4.4.2 Detector Electronics

The CCD detector electronics will be identical for all 32 spectrograph channels. The readout electronic will allow the user to select a fast readout speed (> 1MHz) or a slower speed with lower read noise (<2 e-). All controllers will be running synchronously so that all the CCDs are read together. If TMT decides on using a unique controller design and model for all instruments, HROS’s requirements should easily be met within the chosen design frame. The CCD electronics will most likely consist of:

• computer with auxiliary boards for sequence timing, data capture and network connection to other computers

• detector head electronics with clock drivers, analog bias generation and video signal processing,

• a CCD detector head with pre-amplifier circuitry, • a housekeeping box for temperature control, vacuum sensing and shutter control, • a DC power supply module to generate the 5, ±15, 24 and 30 volts DC power

needed by the detector head electronics and the housekeeping box

4.5 SOFTWARE DESIGN

The software will allow the user to command, control and monitor the various components of HROS. It will also communicate with the TCS to get and return any relevant information related to the operation of the telescope and instrument (guiding offset, WFS, current pointing for the ADC and image rotator).

4.5.1 Instrument Control Software

The Instrument Control Software (ICS) will be responsible for receiving commands from the user or the TCC, controlling and monitoring the various systems that make up the instrument. This software will be designed within the constraints of the TMT instrument control requirements and communication standards to be developed. Standard software engineering methods (use cases, UML, etc.) will be used during the design and development process of the ICS.



4.5.2 Data Flow and Handling

The instrument’s data flow will consist primarily of the image download from each detector but also the commanding, monitoring and real-time feedback to the TCS from the WFS, and auto-guider. Each CCD controller will be linked to the ICC via fiber optics cable to a PCI-express card. The high speed data link coupled with the high speed data bus should provide for a read time of ~10 seconds for all 32 detectors assuming a 4 channel readout per detector. Once the data in read and stored in the computer’s memory, the system will then write out properly formatted FITS files on large capacity hard drives. The commanding and monitoring of the various parts of the instrument as well as the communication with the TCS will all be handled by the ICC. We will evaluate and select the best approach to send the information from the AO system and the auto-guider back to the TCS according to the TMT requirements.

4.5.3 Observation Planning Tools

A tool will be provided to plan observations. This tool will allow the users to select any of the instrument modes (on-chip binning, grating setting if any), the type of object to be observed, and an estimate of magnitude as well as expected seeing conditions and exposure time. The output will be a full spectrum covering the whole range of the spectrograph using the measured efficiencies of the various optical components and the wavelength quantum efficiencies of each CCD and dichroics. The user will be able to easily determine the exposure time required to obtained the desired signal-to-noise.

4.5.4 Data Analysis Tools

Data analysis tools will be provided to the user community to specifically handle the required analysis tasks for HROS data. These will most likely be written in either IRAF, IDL or Python. Guidelines will be provided in the CU-HROS User’s Manual on the type and amount of calibration data required to successfully use the analysis tools and to extract fully calibrated spectrum. The standard set of CCD calibration tools will be provided: bias, dark field subtraction, flat fielding, bad pixel map, cosmic ray removal, etc. Additional spectroscopic tools will also be provided to extract the spectra, subtract sky features, combine all 32 spectra together and correct the wavelength regions where the dichroics overlap. The tools will provide a way of processing the data in a batch mode as well as doing each step manually with the option for the user to inspect the data at every step.



4.6 KEY AND HIGH RISK COMPONENTS

4.6.1 FIFU

This is recognized as a high risk component, both because the technology is relatively young, and because of the importance of a high performance fiber IFU to CU-HROS performance. In order to mitigate risk, we will also explore fallback options, chief amongst these being a conventional image slicer. The slicer will not allow for concentration of core energy like a spiral packed fiber IFU, nor will it enable multi-target operations. Image slicers do, however, promise somewhat higher efficiency.

4.6.1.1 Manufacturing Risks

Key to a broadband fiber IFU is a high efficiency microlens array, including, if possible, a broadband antireflection coating. Several manufacturers can produce arbitrary arrays in a fused silica substrate, and samples will be obtained in the design study phase from multiple manufacturers and examined for uniformity, fill factor, transmission and microroughness. Another risk is our ability to coregister the fiber bundle to the microlens array, although this has not been a significant problem for much higher density arrays.

4.6.1.2 Cost Multipliers

High performance microlens arrays represent a potential cost multiplier. However, the component in question are required in small quantities, and are relatively inexpensive when compared to other systems, such as the camera optics, dichroic mirrors, or detectors, so even a large increase in expense of this element would have little impact on the overall budget.

4.6.2 Dichroic Tree

The dichroic mirror array represents another high risk area because performance of the CU-HROS depends on the acquisition of large format (200 – 250 mm clear aperture) dichroic mirrors with uniformly high efficiency and sharp transition edges. This design would have been out of the question until recently, but advances in the field have radically increased the quality of optics available. We have identified three manufacturers capable of producing the dichroic mirrors for CU-HROS: Research Electro-Optics, ZC&R Coatings, and Barr Associates. We have data on actual optics from Barr Associates, and so have baselined them as the vendor, although we would obtain and compare test articles from all three manufacturers before committing to one.



4.6.2.1 Throughput

We require at least 95% transmission/reflection and sharp transmission/reflection edges (5nm transition from 90% transmission to 90% reflection). Barr Associates has achieved this in the past and they express confidence in their ability to achieve this for our project.


We mitigate manufacturing risk in two ways: First, we will work with multiple manufacturers to obtain test articles prior to instrument build, and second, we have made every effort to tailor our requirements to the capabilities of the manufacturers. For example, we only require a single transition, and are not constraining performance outside of the required operating wavelengths. This simplifies manufacture, especially on progressively deeper banks of optics, as the mirrors only need to perform well over a narrow wavelength range.


If larger optics are found to substantially reduce light loss from vignetting, then this will represent a substantial impact on cost. However, the dichroic mirrors are substantially less expensive than other components such as the camera optics or the detectors, and a significant (2x) increase in cost would raise the total instrument cost by less than 5%.

4.6.3 Cameras

The performance of the spectrograph cameras is crucial to the success of CU-HROS. Designs exist that would produce adequate imaging across the required spectral bandpass, but the challenge is finding cost-effective solutions that have adequate performance. We will consider refractive, reflective and hybrid camera systems. We would prefer all-reflective designs as one design can be used for all channels. Refractive designs exist and may be the most suitable design choice.


The advantage to building many particular units is amorization of the design, engineering, and test setup costs. Fewer required designs for the camera (~3 total) will result in significant cost savings for testing and processing.

4.6.4 Detectors

We will require custom detectors for CU-HROS. These detectors, while large, will not be the largest detector units made by either of the two major vendors we have contacted (Fairchild and e2v). The technology to form chips on the desired wafer size already



exists and we will capitalize on their development for other projects. Similarly, pixel size and well depth do not present a challenge to these vendors.


The large format increases risk, as the yield will decrease with only one chip per wafer. Therefore more wafer runs may be required to complete our fill order. Fairchild and e2v are both currently working to increase yields and produce very large format monolith CCDs, and their development work is likely to continue through the future.


The detectors will take significant advantage of multiple units. By using the same format, the set up and design cost will be amortized over all produced units. It would be less expensive to use devices that exist for other purposes but at present no suitable devices exist.

4.7 INSTRUMENT MAINTENANCE

4.7.1 Maintenance Requirements

The baseline design for CU-HROS features very few mechanisms, reducing maintenance requirements. Some design alternatives, especially those involving multi-resolution operations, do involve numerous mechanisms, but these are relatively simple, reliable, and highly duplicated, so a limited supply of spare parts would support continuous instrument operations.

4.7.2 Spare Parts List

A list of recommended spare parts for CU-HROS is given in Table 9.

CU-HROS On-Site Spares System Qty of Spares Down Time Criticality ADC motor and bearings

1 set Single day shift 3

Derotator motor and bearings


AO feed mirror positioning mechanism1


AO trombone1 1 set of mechanical parts Single day shift 2



Fast Steering Mirror1

1 assembly Single day shift 3

Grating select mechanism2

2 motors and one stage Single day shift3 3

CCD 2 CCDs: one for blue coverage, one for red

24 hours3 3

CCD electronics 2 boards No downtime if external 3 CCD cooler 2 assemblies No downtime for

controller, 1 day shift for cooling head3

3

Cryo pump4 50% overcapacity on line, one spare compressor and cryopump

No downtime 1

Chamber roughing pump4

1 spare pump assembly No downtime 1

Notes: Criticality codes: 1: Depends on AO functionality 1: No impact on operations 2: Not applicable to baseline design 2: Limited reduction in performance 3: If in vacuum, then 48 hrs downtime likely 3: Some loss of science 4: Not applicable unless a vacuum chamber is used.

4: Single point failure leads instrument shutdown

Table 9: Spare parts list for CU-HROS

4.7.3 Instrument Downtime

The only components of CU-HROS requiring regular maintenance are the cryocoolers for the CCDs and the roughing pumps and cryo-pumps (if required). Periodic servicing of the cryocooler compressors will limit downtime, and cryo-pump overcapacity will allow continued operation even if a single pump fails. If the system is not under vacuum, then all repairs can be carried out with little difficulty during daytime operations. If the system is under vacuum, then we anticipate a minimum of 48 hours downtime if the vacuum chamber is opened. This is dominated by thermal stabilization time after pumpdown.

5. PERFORMANCE ESTIMATES

5.1 SENSITIVITY CALCULATOR



We have written a sensitivity calculator to predict CU-HROS performance. We first take the system throughput in conjunction with the telescope throughput to estimate the signal from a source on the sky. We then compute the contribution of sky background and read noise into signal to noise of the observation. The performance for CU-HROS will be highly even across the bandpass due to individual optimization. The components with significant efficiency dependence on wavelength (grating, CCDs) can be optimized for each channel. There are effects due to wavelength that we are not considering, but these have minor effects on relative throughputs (e.g., the throughput of fused silica has a wavelength dependence that we are not accounting for due to its minor impact on the system). Sky values similar to Mauna Kea (taken from the TMT Greenbook) were used to compute the contribution of background to the spectrum. No attempt was made to fit individual OH lines or other atmospheric features. The read noise was assumed to be similar to existing CCD arrays. The e2v 44-82 2k by 4k back illuminated device (similar to the device we would use) has a specification of 2.5 electrons rms per read.

2)()(

)(

RNtBtS

tS

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+!+!

!=

In order to appropriately compute the contribution of the background and read noise, we include the spread of light from each resolution element over a finite number of pixels. With the fiber fed unit (FIFU) deployed, the spot is sampled (at R = 100,000) at 0.1 by 0.1 arcsecond pitch with 2 pixels per 0.1 by 0.1 arcsecond at the final camera. With good seeing (~0.5 arcseconds), the light will be spread out over ~ 40 pixels at the 90% encircled energy level. While this is a large number of pixels and read noise is a significant noise source, the number of pixels is unavoidable with large telescopes and arcsecond scale images. The two dominant places to improve HROS faint object performance are in the spot size on the FIFU — a smaller spot size spreads light over fewer pixels therefore with lower background and read noise for a given observation — and reducing the read noise on the detectors. In broad performance, HROS will produce a signal to noise spectrum of 100 (minimum per spectral resolution element) at 100,000 resolution (minimum) of an object with mAB equal to 16.5 in six hours with 1.0 arcsecond 90% encircled energy. Signal to noise is a fairly flat function of wavelength (varies from 100 to 101 due to background). For a calibration object, a magnitude (AB) 4.0 object will produce 75,000 counts per pixel with 0.5 arcsecond seeing in one second. More typically, we would recommend a much



fainter object (m ~ 8.5) to eliminate the chance of saturating many pixels over the detector.

Resolution Seeing (90% encircled) Moon Phase Limiting mAB at S/N of 100 100,000 1.0 Dark 17.5 100,000 1.0 Bright 17.5 100,000 0.5 Dark 18.9 100,000 0.5 Bright 18.9 50,000 0.5 Dark 19.4 20,000 0.5 Dark 19.5

Table 10: Limiting Magnitude for 6-hour total exposures with 5 reads, with on-chip binning employed for the lower resolution observations.

Note that background (i.e., bright time) is a minor effect – the noise characteristics are dominated by read noise at all wavelengths for the CU-HROS design.

5.2 TECHNICAL VIABILITY OF SCIENCE PROGRAMS

With an overall throughput of ~20%, the CU-HROS design maintains the aperture advantage of TMT over current telescopes when compared to the throughputs of current cross-dispersed designs at their blaze wavelengths. Because the CU-HROS design maintains this throughput over the full optical waveband, our concept improves upon the aperture advantage of TMT at most wavelengths compared to current instruments, which show rapid drop-offs in throughput away from the blaze. Based on past delivery histories, we also believe that some of the component throughputs, in particular for the dichroics and the gratings, will be overachieved by the manufacturers, resulting in overall improved instrument sensitivities. Based on the instrument throughput numbers given above, the CU-HROS concept will support R=100,000 spectroscopy of the full optical waveband in a single exposure at S/N=100 for objects as faint as V~20, and R=20,000 spectroscopy at S/N = 20 to V<22. This is sufficient to support the science goals for HROS, as outlined in the TMT DSC and the CU-HROS IOCDD. In particular, this sensitivity will allow HROS to obtain background QSOs for IGM studies to a spatial sampling of 1 QSO per ~7 square arcminutes and to pick up a large sample of gamma-ray bursts as distant probes (where the mean GRB magnitude from Swift detections is 21.5 at 12 hours after the burst). It will also enable HROS to probe the red giant branches of several dwarf irregular galaxies in the Local Group, observe thousands of very metal poor stars in the Galactic halo, and more than triple the number of M stars available for planet searches.

6. TRADE STUDIES



6.1 FOREOPTICS

6.1.1 Atmospheric Dispersion Compensation

A standard ADC design with fused silica prisms will likely cost ~$100K. We would consider alternative design but with the confined zenith angle of TMT, the ADC is a low risk component. Decreasing the field of view to 10 arcsec may impact performance by decreasing the number of natural guide stars available to CU-HROS. Based on existing designs, the smaller field of view ADC would cost $40K.

6.1.2 Acquisition and Guiding

The acquisition and guiding camera may be be incorporated into a single unit at the cost of reducing system throughput (by using light from a pupil to do both all the time). Alternately, a single camera at the reimaged field could be used if TMT pointing and initial acquisition are accurate enough to locate the guide star on the field of view. Without further specifications for TMT performance, we have difficulty analyzing the advantages of different solution paths at this point.

6.1.3 Calibration

Subjects to consider include the need for multiple wavelength calibration sources, the best means for injecting the calibration signal into the instrument, and the selection of appropriate photometric calibration sources. Because a reliable and repeatable wavelength solution is essential for high resolution spectroscopy, the baseline design calls for three separate wavelength calibration sources: an iodine absorption cell, a set of emission line lamps, and a laser frequency comb. The stability, reliability and cost effectiveness of these sources will be evaluated. Internal quartz lamps should be available for flat-fielding, but HROS should also contain a diffuse filter that enables the acquisition of daytime sky flats to support sufficient flat-field S/N at blue wavelengths, where lamps are less efficient.

6.2 ADAPTIVE OPTICS

We will need to balance the gains of AO versus the cost of implementation. If there is need to accommodate multiple laser guide stars or a single guide star that will alter the foreoptics design, we will have to consider the impact of optical laser light falling into the spectrograph aperture. The CU-HROS design uses a fiber IFU to restructure the point spread function (PSF) of TMT into a radially concentrated pseudo-slit. As a result, decreasing the instrument PSF



will not directly impact CU-HROS spectral resolution, but it will spread the light over less of the pseudo-slit in the cross dispersion direction, making on-chip binning more effective and decreasing exposure times required to reach a given target. So while CU-HROS does not rely on AO for its spectral resolution, it would still benefit from a high performance AO system. One major difficulty with both approaches is the impact of putting an extremely bright signal onto the CU-HROS entrance slit right in the center of the instrument wavelength band (589 nm). This signal is many orders of magnitude brighter than the astronomical signal, and would probably overwhelm the 589 nm channel, as well creating a very bright source of scattered light. Another difficulty is that CU-HROS operates beyond the blue limit of the required operating ranges for the AO systems (TMT.PSC.05.001.REL15), and little data or modeling is available for short wavelength performance (although see the discussion of models for HROS in Section 4.2.2.3). Further studies in the design phase for this section would compare in greater detail the cost, benefits and impacts of GLAO and SLGLAO against a system with no AO, and, by extension, against one another. This trade study is not as critical for the instrument design as others. While it will impact the foreoptics, it does not have any schedule impact on the principal long lead components (IFU, collimator, dichroics or spectrograph benches).

6.2.1 SLGLAO

SLGLAO should perform nearly as well over the very narrow field of view of HROS as would GLAO, but without requiring the larger field of view foreoptics demanded by the full constellation of laser guide stars (70 arcsec for SLGLAO vs. approximately 210 arcsec for GLAO). However, the reduced field of view makes it somewhat challenging to find a bright natural guide star in the instrument field of regard. Note that if the CU-HROS were to switch to an image slicer rather than the FIFU, then there is less advantage to using AO, since we will not be able to create a radially concentrated output with an image slicer’s pseudo-slit.

6.2.2 GLAO

While SLGLAO has the advantage of requiring smaller foreoptics than GLAO, there may be advantages to working with the wider field of view demand by the full constellation of LGS for GLAO: First off, with the full constellation surrounding the critical central region of the focal plane, rather than being centered on the CU-HROS field of regard,



there will be less 589 nm light entering the CU-HROS instrument cavity, and so less scattered light. It might be possible to eliminate this with additional dichroic mirrors to protect the 589 nm channel of HROS, or, failing that, to shutter this channel but continue operation in other channels. Additionally, with the wider field of regard for the foreoptics, we increase the probability of locating a natural guide star. This trade study would balance the additional costs of much larger foreoptics and multiple wavefront sensor optical systems against the advantages promised by AO, and against the increased scattered light and loss of science in the 589 channel. As with SLGLAO, if we opt for an image slicer rather than a fiber IFU, it is not clear that GLAO would provide any significant benefit.

6.3 OPTIONS FOR IMAGE SLICING

Our ability to produce 100,000 resolution relies on dissecting the image to 0.1 arcsec wide sub-images. Without that capability, the resolution would become slit limited and either present lower resolution for poor seeing or present a significant decrease in throughput at the slit. We do not consider CU-HROS to be a viable instrument without some form of image slicing. The two principal options available are conventional image slicers and fiber fed IFUs. Both systems could be used to create a 0.1 arc second pseudo-slit. The instrument described in this feasibility study assumes a fiber optic IFU, but that is a preference rather than an absolute requirement. In Table 11, we compare advantages and disadvantages of the two options. On balance, we consider the advantages of the FIFU to outweigh the minor benefits associated with adopting an image slicer. This question would be re-examined in greater detail during a design study, however. If CU-HROS were to switch to an image slicer, the difference in output f/# would demand that this trade be completed early in the design process, since it directly impacts the collimator design, and, therefore, all downstream optics. Fiber IFU vs. Image Slicer Considerations Advantages Disadvantages

Allows radial concentration of telescope PSF, increasing S/N.

Fast output beam complicates collimator design.

Fiber IFU

Can be moved if multiobject capability is desired.

Inherently less efficient than the 2 or 3 bounces of an image slicer.



Very easy to inject calibration signal.

Can be used to extract partial signal for fast centroiding.

Supports spectroscopy of extended objects

Relatively higher efficiency Does not concentrate central region of beam on central region of detector

Image Slicer

Slow output beam simplifies downstream optics

Cannot be repositioned without additional relay optics

Table 11: Comparing fiber IFUs to Image Slicers

6.4 OBSERVING MODES AND EFFICIENCY

Note that in its default configuration, CU-HROS has only one observing mode. We deliberately adopted a design that would be simple and flexible to allow HROS to function quickly and easily under a variety of observing conditions. As a result, trade studies for observing modes and observing efficiency are limited. In Section 8, we do evaluate descope options driven by cost that would contain multiple observing settings; here, we discuss the effects these descopes will have on observations.

6.4.1 Spatial Sampling

The spatial sampling in the FIFU (0.1 arcsec per resolution element or 0.05 arcsec per pixel) was chosen to minimize the optics sizes of the instrument as well as to maintain Nyquist sampling in the event of AO-enabled small image sizes. All of the science cases outlined in the IOCDD are for point source spectroscopy, but the presence of an IFU will support observations under good seeing conditions and/or the deployment of the IFUs over the extended object. The number of IFUs was driven by a combination of available space in the spatial direction on our CCDs and the need to sample object plus sky. There are cost savings (particularly in development costs) in keeping the IFUs in a fixed configuration. The advantages are allowing the IFUs to move are twofold: enabling multiplexing of target fields with sufficient population density that multiple targets appear in our relatively small field of view (70 arcsec), and allowing the selection of regions of truly blank sky for sky sampling in dense regions.

6.4.2 Sensitivity



The throughput of CU-HROS will have variations of only a few percent (e.g. 17.5 % at 310 nm rising to 19% at 400 nm) across the entire wavelength range. By using each channel at peak, we expect throughput at the 19% level modulo the presence of the dichroic gaps. Depending on the performance of deep depletion CCDs, there may be some performance decrease at very long wavelengths due to silicon becoming transparent. See 4.2.3.2 for a more complete discussion.

6.4.3 Simultaneous Wavelength Range

The requirement was for a spectrograph that will function from 0.3 microns to 1.1 microns. We felt that complete simultaneous wavelength coverage is a driver for some of the scientific cases. If we descope the design such that two grating tilt settings are needed to paint the full optical waveband, the result is a direct loss in observing efficiency for any observation interested in more than one narrow wavelength region. The descope design (the use of 16 channels with two-position gratings) results in a single exposure covering half the optical waveband in non-contiguous “chunks”. This pattern is unusual for astronomical instruments, but has been employed before (for example, in the NUV channel of the Cosmic Origins Spectrograph for Hubble) and represents an easy way to decrease instrument costs without complicating design and operations.

6.4.4 Spectral Resolution

We would have the system set at 100,000 resolution with only the shutter as a moving part. This way if the observer desires lower resolution, the logical choice will be to perform on-chip binning to reduce the impact of read noise. We have presented another descope option, to reduce the maximum spectral resolution to R=60,000 or R=80,000. We do not support these descopes because they will result in an instrument that will be extremely difficult to retrofit at a later date for higher resolution and because we strongly believe that one of the strengths of HROS on the TMT will be the ability to obtain the combination of extremely high quality spectra at spectral resolutions unavailable to current telescopes.

6.5 DEPLOYMENT ON NASMYTH PLATFORM

CU-HROS is compact enough to be accommodated on the Nasmyth platform with little difficulty. If necessary, we can reduce our instrument footprint by increasing the degree of stacking, although we would then need to increase the instrument height. This can be accomplished either by growing the instrument up from the existing platform, or by penetrating below the Nasmyth platform. The fiber IFU decouples the instrument performance from the telescope beam speed, so that changes in the telescope beam speed and back focal distance can be accommodated in the foreoptics.



6.6 VACUUM CHAMBER

The vacuum chamber is considered the default design because of tight tolerances on dichroic mirror acoustic frequency vibrations and on component and bench level thermal gradients. These are discussed in section 8.6. Once a final instrument configuration is adopted, detailed optomechanical and thermal analysis of the instrument will be carried out, and the decision to retain or descope the vacuum chamber for the instrument cavity will be made.

7. UPGRADE OPTIONS

7.1 J-BAND ARM

By ensuring the dichroics have appropriate performance, we can add 100,000 resolution spectroscopy across the J-band (from 110 to 130 nm) with three channels. The extra channels would be added to the area around the red-most channel in the current design (#32) and would have no effect on the rest of CU-HROS. The channels themselves may present issues, in particular with the requirement to change detector technology to HgCdTe devices. The overall impact would be minimal as the spectrograph bench interface to the controllers can be configured identically to the existing CCD controllers.

7.2 MULTI-OBJECT CAPABILITY

The default design uses a fixed pattern for the FIFUs with a clear upgrade path to positionable units. Elimination of the positioning mechanism for the FIFUs still allows for sky sampling by placing a second and third FIFU bundle in a fixed position away from the central lenslet array, but would make acquiring multiple targets unlikely. The central array could be enlarged, either extended linearly in the hope of acquiring a second target or sampling an extended object, or extended radially, allowing us to sample more of the telescope PSF in poor seeing conditions. We have examined the costs associated with making the FIFUs positionable from the start. For positioning the FIFUs for multi-target programs, we have based our concept on the one used on Hectochelle on the MMT. The position of the fibers needs to have similar accuracy to the Hectochelle design (100 microns) and will require many fewer mechanisms (instead of 300 fibers, ~5 fiber bundles would be used). The system would use extending arms instead of robot positioners. A magnetic lock to the focal surface would be the preferred locking mechanism. Also, the guide stars could be accessed at this point with ease. Estimated cost for positionable FIFUs is ~$100,000. However, the benefit to mechanizing the fiber bundles is not clear without a larger field of view: while there are some science applications that could benefit from multiplexing with the 70"



FOV (lensed QSOs, extragalactic stellar spectroscopy), most require a larger FOV to take advantage of multiplexing. Implementing a larger field of view has downstream consequences for cost that would have to be taken into account for large-field multiplexing (ADC, derotator, reimaging optics, etc).

7.3 AO INTERFACE: SLGLAO TO GLAO

Changing the AO design that CU-HROS can accommodate from SLGLAO to GLAO represents a substantial cost increase: the foreoptic field of view increases from 70 arcsec to approximately 135 arcsec, and it would require the addition of 4 or 5 wavefront sensors beyond the single system required for SLGLAO. It is unclear if there would be substantially better performance obtained with GLAO over SLGLAO for a point source spectrograph, but the annular array of laser guide stars may be preferable from the perspective of stray and scattered light.

7.4 R>100,000 SPECTROSCOPY

We can consider adding an R > 100,000 mode via careful dithering of the fiber outputs (and therefore the spectrum on the chip) to push the resolution closer to single pixel levels. The dithering technique would be similar to the one used on Hubble Space Telescope Advanced Camera for Surveys Wide Field Channel. Alternate concepts include smaller pixel pitch which will negatively impact read time, read noise, and well depth. There will be similar issues in supporting lower resolution observations as in the conventional design (increased number of pixel illumination, therefore increasing the read noise contribution).

8. DESCOPE OPTIONS

8.1 R≤60000 ONLY

One descope option for HROS is to reduce the default resolution from 100,000 to 60,000. This will reduce the number of channels from 32 to 19. A disadvantage to descoping to R of 60,000 is the difficulty in recovering higher resolution. Solutions to reacquire the full resolution exist, but have impacts on observing efficiency. A possible recover option would be to use tippable gratings instead of fixed pieces. With the precision needed, this would add 19 (baseline configuration) tip mechanisms for each of the 19 channels. By tipping the grating and keeping

!

" #$ constant, we can tip the grating for a new bandpass in each channel. There may be impacts to the efficiency due to a larger bandpass for each grating, but until detailed models of the grating profiles are made, we cannot discuss this in other than broad terms.



Cost to implement the 100,000 resolution rescope is mitigated by the interleaved spectra, which will significantly reduce the precision required to position the gratings. As the wavelength calibration falls in every spectra, the actual position of the grating is easily determined without requiring high absolute precision mechanisms.

8.2 R=100,000 THROUGH NON-CONTIGUOUS WAVELENGTH COVERAGE

CU-HROS could achieve full resolution with half the spectrograph benches and dichroic mirrors, but at the expense of not being able to obtain the full spectrum in a single observation. The dichroic tree would consist of only 15 units in 4 banks (somewhat increasing efficiency), so the spectral grasp of each bench would be twice as large. Each spectrometer bench would be fitted with two gratings: one for R=50,000 operations, and a second for R=100,000 operations. In the R=50,000 mode, the gratings would remain fixed and the entire spectrum could be obtained in a single observation. In the R=100,000 mode, the high resolution grating would be rotated into first one position, and then to a second orientation. Each grating position would allow for slightly more than half of the spectrum to be observed, but in non-continuous bands, with the second grating position being needed to fill in the gaps. This option has the advantage of saving roughly $2.8M, but at the expense of significantly reducing observing efficiency. However, if the lower resolution instrument is seen as the baseline, this represents an relatively inexpensive ($0.8M) incremental upgrade if high resolution is deemed sufficiently important in the future.

8.3 BLUE CUTOFF AT 340 NM

Blue cutoff due to telescope optical coatings will remove four channels from the blue end of the spectrograph. The technical impact is minimal. In principle, the dichroics can be arranged to have blue channels added as funding becomes available with minimum impact to the rest of CU-HROS.

8.4 SINGLE OBJECT SPECTROSCOPY ONLY

The 70" field of view was chosen to accommodate acquisition of a guide star and future AO capability through SLGLAO or laser tomography. It is possible to achieve CU-HROS with the original default 10" FOV. If light from a pickoff at the pupil is used, CU-HROS could guide on the target directly but at the cost of decreased throughput, which is not in keeping with the goal of maximizing sensitivity for HROS. Note, also, that some AO capability may be required of all instruments when the deformable secondary mirror is implemented on TMT.



8.5 ELIMINATION OF LASER FREQUENCY COMB

While this technology represents a profound advance in spectroscopic calibration capability, it is also undemonstrated and quite expensive. Implementation would provide uniformly distributed wavelength reference throughout the instrument bandpass, but this can be accomplished to a lesser extent using established, less costly technologies, making this a prime descope target. This technology represents an area where it would be useful to keep an eye on advances in the field for possible future use.

8.6 ELIMINATION OF VACUUM ENCLOSURE FOR OPTICS CAVITY

Preliminary tolerance analysis indicates that the two principal threats to CU-HROS performance are vibration of the dichroic mirrors (see Section 4.2.2.7.2) and thermal motion of the detectors with respect to the grating bench. Our preferred solution is to place the bulk of the instrument in a vacuum chamber to mitigate acoustical and thermal vibration effects. The vacuum chamber will not need to be operated at particularly low pressures (<10-5 torr) to achieve the required effects. Cycling time would therefore be as short as a few hours to bring the enclosure to air and to pump back to vacuum, and up to 48 hours to reach a stable vacuum pressure. The vacuum chamber is heavy, however, and may be more expensive than the creation of a clean, thermally controlled room. We would therefore revisit this question in the design study phase to determine if the vacuum chamber is the best choice for the CU-HROS enclosure.

9. COST ESTIMATES

9.1 DEFAULT DESIGN

Estimated cost by subsystem is summarized in Table 12.

CU-HROS Cost Summary System Description Cost Estimate 1 Foreoptics + Calibration Systems $ 0.82 M 2 Fiber IFU $ 0.36 M 3 Collimator $ 0.15 M 4 Dichroic Mirrors (31x) $ 0.64 M 5 Gratings (32x, holographic) $ 1.28 M 6 Cameras $ 3.20 M 7 Detectors $ 4.02 M 8 Mechanical $ 2.00 M 9 Personnel (54 man-years) $ 8.47 M Full system cost $ 21.12 M



Full system cost with 1.3 safety factor: $ 27.20 M Table 12: Cost Estimates for the Default CU-HROS Concept

9.1.1 Budget Narrative

Line 1, “foreoptics and calibration systems,” is dominated by the laser frequency comb and by the reimaging optics. Based on existing systems at NIST in Boulder, Colorado, the expense for the laser frequency comb is estimated to be $ 250,000 for a single system, with and additional $ 100,000 budgeted for 20:1 frequency filtering, tailored attenuators, and preparation for use in the field. Estimated costs for the derotator will be approximately $50K, depending on the size of the field of view it will be required to support. However, since this is a series of flat mirrors with mechanisms, the cost of the mirrors is much less sensitive to the size of the substrate than the ADC. The cost for the reimaging optics are drivien byt the field of vew, but need not brow beyond 70 arc seconds in the event that GLAO is selected since GLAO pickoff occures before the reimaging optics. Line 2, “Fiber IFUs,” is dominated by the micropositioning system. As with all hardware estimates, this does not include labor costs, which are covered in Line 9. Line 3, “Collimator,” is based on estimates from previous instruments and vendors for the desired optics sizes. Line 4, “Dichroic mirrors,” is from a cost estimate provided by Barr Associates. This estimate came in between other estimates; e.g., by Research ElectroOptics, Inc. However, Barr Associates had considerable experience in large format, high quality dichroic mirrors, so their estimate is used. Line 5, “Diffraction gratings,” is based on a quote provided by Jobin Yvon. We have extensive experience with this vendor and have consistently received accurate cost estimates and excellent optics. Line 6, “Camera optics for the spectrograph benches,” is based on initial designs provided to Goodrich Corporation. Line 7, “Detectors,” is based on estimates provided by Fairchild and e2V. Ordering 32 identical CCD substrates results in substantial cost savings over the order of a single, large CCD mosaic. Line 8, “Mechanical Systems,” is dominated by the vacuum chamber and the clean room. Costs for these are based on scaling the facilities at CASA in volume and applying



inflation for the 10 years since these facilities were constructed. We operate a 90m3 vacuum chamber in a 600m3 class 1000 clean room. We anticipate requiring a 140m3 vacuum chamber and a 400m3 class 10,000 clean room. Initial costs for the CASA systems were $400,000 for the clean room and $400,000 for the vacuum chamber. Line 9, “Personnel,” is based on the development schedule discussed in section 10.2. In general, cost data provided include spares as outlined in Table 9.

9.2 UPGRADE COSTS

Upgrade Cost Summary Descope Option Cost Delta Discussion

Addition of J-Arm

+$1,366 K

Addition of at least 3 spectrometer benches and 3 dichroics

SLGLAO +$253 K Addition of WFS and associated hardware GLAO +$1,618

K (estimate)

Increased FOV foreoptics, additional WFS and feed optics

Table 13: Upgrade Costs for CU-HROS

9.3 DESCOPE COST SAVINGS

Descope Cost Savings Summary Descope Option Cost Delta Discussion

Small (10") FOV -$ 254 K Eliminates need for 70” FOV foreoptics R=60,000 -$ 3,895 K Eliminates 40% of spectrometers, dichroics. R=50,000/100,000 -$ 3,829 K Eliminates 50% of spectrometers, dichroics.

Requires same number of gratings, adds 16 mechanisms.

340mn cutoff -$ 1,823

Eliminates 4 spectrometer benches and 2 dichroics

Single object spectroscopy

-$ 65 K Reduces mechanism count

No LFC calibration source

-$ 455 K Reduction in wavelength calibration capability

Non-vacuum enclosure for optical cavity

-$ 1,137 K Vacuum chamber will be implemented only if deemed necessary for instrument stability



Table 14: Descope Cost Savings for CU-HROS

10. INSTRUMENT DEVELOPMENT PROGRAM

This section presents a top-level work breakdown structure (WBS) outlining all the tasks necessary to design, fabricate, integrate, test, and commission HROS and their associated costs.

10.1 WORK BREAKDOWN STRUCTURE

CU-HROS Work Breakdown Structure WBS Id. Task Description 1 HROS Design/Definition

1.1 AO trade study 1.2 Resolution trade study 1.3 FIFU trade study 1.4 Multi-object trade study 1.5 Optomechanical tolerance and vacuum chamber evaluation 1.6 Upgrade/descope paths 1.7 Focal plane and foreoptics design definition

2 HROS Design 2.1 HROS Foreoptics 2.2 HROS Core Optics 2.3 Structural/Thermal Design 2.4 Electrical 2.5 Instrument interface software

3 HROS Fabrication 3.1 Optical Systems Procurement 3.2 Bonding and Mounting 3.3 Component testing and verification 3.4 Subsystem Assembly 3.5 HROS Integration

4 HROS Characterization 4.1 Dispersion 4.2 Channel spectral overlap 4.3 Efficiency 4.4 Stray light, ghosting, etc

5 HROS Data Pipeline Development 5.1 Data product definition 5.2 Detector output characterization 5.3 User Interface Design



5.4 Software build with nominal performance 5.5 Update software for actual performance

Table 15: Top Level WBS for CU-HROS Construction

10.2 DEVELOPMENT SCHEDULE

CU-HROS development will begin with an extended (10 month) project design definition phase. Some design and analysis work will be carried out in parallel as the design matures. Fabrication of foreoptics, collimator and instrument enclosure will be emphasized early in the fabrication phase because of the modular and highly replicated nature of the optics cavity. Fabrication of the enclosure and vacuum chamber should begin as soon as possible at the telescope site in order to accommodate dichroic array components and spectrograph benches as they become available. Foreoptics should be installed at the telescope site as soon as the vacuum enclosure and clean room are completed. As soon as dichroic mirrors and spectrograph benches are available, they can be transported to the telescope site and installed in the instrument cavity. Assuming a January 1, 2008 start date, the build schedule at this point would be constraind by the completion of the Nasmyth platform. A later start would be constrained by the completion of the foreoptics and the first spectrograph benches Initial alignment may begin as soon as the foreoptics and the first bench and a full collun of dichroics (5 optics) are installed, and can proceed without a functioning telescope. Grating, camera, detector and dichroic procurement will be carefully coordinated so that integration may proceed as seamlessly as possible, with a partially functional instrument on line as early as 24 months after project initiation. Calibration and alignment can be carried out as each bench arrives, and the final spectrograph bench will be installed 51 months after project initiation. The fully calibrated and commissioned will be on line 55 months after project initiation. The CU-HROS instrument is capable of providing full resolution without a functioning AO system, or even a fully assembled primary mirror array. This combination of schedule flexibility and insensitivity to telescope performance makes CU-HROS an excellent choice for a first light instrument on the TMT. The CU-HROS development schedule is illustrated in Figure 10-1.



Figure 10-1: CU-HROS Development Schedule

11. BIBLIOGRAPHY

Jeremy Allington-Smith, Grahm Murray, Robert Content, George Dodsworth, Roger Davies, Bryan Miller, Inger Jorensen, Isobel Hook, David Crampton and Richard Murowindsk, “Integral Field Spectroscopy with the Gemini Multiobject Spectrograph. I. Design, Construction, and Testing”, Pub. Astro Soc. Pac. 114, 892-912 (2002). Geraldo Avila, Isabelle Guinouard, Lourent Jocou, Fabien Guillon, Fabrice Balsamo , “Optical fiber link between OzPoz, GIARAFFE and UVES (FLAMES Project)”, Proc SPIE 4841, 997-1007 (2003). Arthur Cox, ed., Allen’s Astrophysical Quantities, 4th edition, Springer, NY (1999) Hans Dekker, Poule E. Nissen, Andreas Kaufer, Franscesca Primas, Sandro D'Odorico, Reinhard Hanuschik , “High S/N, high resolution Image Slicer observations with UVES”, Proc SPIE 4842, 139-150 (2003). Scott A. Diddams, Personal communication, 29 December, 2005 Scott A. Diddams, David J. Jones, Jun Ye, Steve T. Cindif, John L. Hall, Jinendra K. Ranka, Robert S. Windeler, Ronald Holwarth, Thomas Udem and T.W. Hänsch, “Direct Link between Microwave and Optical Frequencies with a 300THz femptosecond Laser Comb”, Phys Rev Letters,Vol 84, No. 22, 5102-5105 (2000). Roger Haynes, Ivan Baldry, Kieth Taylor, David Lee, “Characterisation of cooled ifrared fibers for the Gemini IRMOS”, Proc SPIE 4008, 1203-1214 (2000).



Matthew Alexander Kenworthy, The Development of New Techniques for Integral Field Spectroscopy in Astronomy, Thesis prepared for the Institute of Astronomy and Clare College, Cambridge (1998). Matthew A. Kenworthy, Ian Parry and Keith Taylor , “Spiral Phase A: A Prototype Integral Field Spectrograph for the Anglo-Australian Telescope”, Pub. Astro Soc. Pac. 113, 215-226 (2001). D. Korsch, Reflective Optics, Academic Press, Inc. San Diego, CA (1981) Daniel Schroeder, Astronomical Optics, Academic Press, Inc., San Diego, CA (1987) Smith, et al. 2002 , Proc SPIE 4669, 172-183 Daigo Tomono, Harald Weisz and Reiner Hoffmann , “Fiber IFU unit for the second generation VLT spectrograph KMOS”, Proc SPIE 4841, 390-397 (2003). Thomas Udem, Ronald Holwarth, and T.W. Hänsch, “Optical Frequency Metrology”, Nature, Vol 416, 233-237 (2002)

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