PRELIMINARY DRAFTCMB S-4: Telescope Design Considerations

September 8, 2016

T. Essinger-Hileman, N. Halverson, S. Hanany, M. D. Niemack, S. Padin, S.Parshley, C. Pryke, T. Suzuki, E. Switzer, K. Thompson, et al.

Institutions

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Contents1 Introduction 1

1.1 Optics benchmarks . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2

2 Current CMB telescope designs and maturity 32.1 Small aperture telescope designs . . . . . . . . . . . . . . . . . . . . . . . . . . . 42.2 Large aperture telescope designs . . . . . . . . . . . . . . . . . . . . . . . . . . . 42.3 Table of current telescope/instrument designs . . . . . . . . . . . . . . . . . . . . 4

3 Design studies for next generation telescopes 53.1 Systematic considerations and the CCAT-prime concept . . . . . . . . . . . . . . . 6

4 Telescope engineering to improve systematics 74.1 Monolithic mirrors . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 74.2 Boresight rotation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 84.3 Shields and baffles . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 8

5 Potential future studies and development areas 10

A Optics designs for current projects 12A.1 Advanced ACTPol . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 12A.2 BICEP3 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 13A.3 CLASS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 14A.4 EBEX . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 15A.5 Keck/Spider . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 16A.6 Piper . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 17A.7 Simons Array . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 18A.8 SPT-3G . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 19

B Recent projects using crossed-Dragone telescopes 20B.1 ABS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 20B.2 QUIET . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 21

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1 IntroductionCHARGE: Summarize the current state of the technology and identify R&D efforts neces-sary to advance it for possible use in CMB-S4. CMB-S4 will likely require a scale-up innumber of elements, frequency coverage, and bandwidth relative to current instruments.Because it is searching for lower magnitude signals, it will also require stronger control ofsystematic uncertainties.

Current landscape

• Existing CMB experiments have a range of telescope sizes (∼ 0.3 to 10 m) and styles (coldrefractors and offset Gregory and crossed Dragone reflectors). See the 2013 review articleby Hanany, Niemack and Page [1] for details, and Appendix A for some examples.

Science drivers for CMB-S4

• CMB-S4 will need some telescopes with high angular resolution for lensing, Neff ,∑mν ,

and Dark Energy, but exactly what resolution vs. frequency? This will set the minimum sizefor the large telescopes.

• How many detector-years vs. frequency vs. resolution will be needed? This will determinethe number of telescopes of a given size, which is important because design trade offs aredifferent for single vs. multiple telescope experiments.

• What limits the low-l end of measurements with large telescopes? This will set the maximumtelescope size for constraining r. It may not be possible to get to a detailed understanding intime, in which case we will have to adopt techniques that are most likely to give the lowestintrinsic systematics, e.g., the use of comoving absorptive shields and boresight rotation,both of which may limit telescope size. Thus far, only small telescopes (e.g., BICEP and theKeck Array) have produced results at l . 100, so it seems likely that CMB-S4 will have atleast some small telescopes.

Telescope designs for CMB-S4

• There are many viable options for CMB-S4 telescope designs. For the most part, the tele-scope design is straightforward engineering.

• Small telescopes for CMB-S4 could be cold refractors like BICEP [2], or a new cold re-flecting telescope design. Small telescopes are inexpensive, so there is less incentive to pushon field of view or multi-color pixels; it is easier, and maybe cheaper, to just build moretelescopes.

• The need for more detectors has pushed existing experiments on large telescopes to designsbased on wide-field cameras with multi-color pixels. This is largely the result of making thebest of a single, existing telescope, but the pixel spacing is correct for only one frequency, so

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the mapping speed for other frequencies is degraded, wide-field cameras tend to have manypixels operating with lower image quality, and multi-color pixels come with an efficiencyhit. If CMB-S4 has several large telescopes, designing from scratch, with realistic efficiencyestimates, might lead to a different optimization with each telescope supporting just 1 or 2bands.

• Large telescopes for CMB-S4 could be crossed Dragone designs, with superior image qualityacross a wide field of view [3], or a larger number of less expensive Gregory telescopes, eachwith smaller field of view [4], [5]. We can make a first attempt at this trade off based onidealized optical designs, but the final choice may require a detailed comparison of mappingspeed vs. cost for designs that include realistic estimates of systematic errors.

Systematic errors

• The push for CMB-S4 is to get more raw sensitivity. Systematic errors will certainly be moreimportant for CMB-S4 than for existing experiments, and may limit the sensitivity.

• Pickup due to scattering and sidelobes is likely to be the biggest problem. CMB-S4 designswill need to consider: scattering from optical surfaces (e,g., reflector panel gaps), stops (tominimize any clipping of the beam), comoving shields, and fixed ground screens.

• Several small telescopes have demonstrated good control of pickup using a comoving cylin-drical absorbing shield, in some cases combined with a fixed conical reflective ground screen.Small telescopes for CMB-S4 will likely follow this experience. Existing large telescopeshave far from ideal shields, largely because of arbitrary notions about what was possible,so we should be able to do a much better job for large CMB-S4 telescopes. The key is toinclude shielding in the initial telescope concept, rather than add it as an afterthought.

• A framework will be needed for estimating the impact of pickup. We might start with a sim-ple beam model, based on some generic description of sidelobes, and investigate the impactof pickup on, e.g., constraints on r. This could be followed up with full beam simulationsfor real antenna and shield designs. The analysis will be challenging for large telescopesbecause the surfaces are huge compared with the wavelength.

Other constraints

• Telescope designs for CMB-S4 must account for operational constraints (e.g., weather, re-dundancy/reliability, and limitations on power consumption) and project constraints (e.g., adown select is pointless if the only way to engage critical partners is by supporting multipledesigns).

1.1 Optics benchmarksDescribe and define important optical quantities.

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2 Current CMB telescope designs and maturityThe current generation of CMB telescope designs incorporate lesssons learned from decades ofmeasurement experience. The nine current-generation experiments and two previous-generationtelescope designs have a wide variety of optical design approaches, including both refractive andreflective primary aperatures, Gregorian and crossed-Dragone mirror configurations, single- andmulti-camera systems (also called tertiary re-imaging optics or optics tubes), and polarizationmodulation mechanisms including half-wave plates (HWP), reflective variable-delay polarizationmodulators (VPM), telescope boresight rotation, sky-rotation, or no polarization modulation.

Small aperture designs, < 1 m diameter, are used to probe for the inflationary gravitationalwave B-mode signal, whereas large aperture designs, > 1m diameter, are used to measure theCMB lensing signal, and > 5m designs can also measure arcminute-scale secondary anisotropies.Small- and large-class telescope designs are discussed separately in the sections below.

Reduction and mitigation of spurious systematic optical signals is a primary design considera-tion for modern CMB telescopes, and has become increasingly important as weaker cosmologicalsignals are being probed. All current designs are either refractive, or off-axis reflective to eliminatefeed legs, which scatter light into far sidelobes, in the optical path. Sources that are chief culpritsfor introducing scan-synchronous signals are the ground, the sun and moon, and the galaxy. Toreduce these signals, reflective ground shields and absorbing baffles, both co-moving with the tele-scope and fixed, are used to reflect scattered light to the cold sky, preferably near the main beam,or absorb it.

Telescope designs for CMB polarization measurements must also minimize spurious polarizedsignals, or polarization systematics, introduced by the telescope, and stabilize those that remainso that they can be characterized and removed during analysis. Instrumental polarization, or thecreation of a spurious polarized signal from unpolarized emission, is caused by assymmetries in theoptics or detectors with respect to the two orthogonal linear polarization vectors. Culprits includemismatches in gain calibration, beam shape, or bandpass in differenced detectors, and by pointingerrors. Cross-polarization, or spurious rotation of the polarization angle, can result from anglemis-calibration and pointing re-construction errors. Off-axis dual reflectors introduce unavoidableinstrumental polarization, but the location and tilt of the secondary with respect to the primary canbe chosen to obey the Mizuguchi-Dragone condition [6, 7], which eliminates cross-polarization atthe field center and also improves the diffraction limited field of view (FOV).

Polarization systematics can also be mitigated by introducing a polarization modulator, whichrotates the polarization angle to which a detector is sensitive. Polarization modulators can onlymitigate systematics that are introduced in the optical path between the modulator and detector.

Finally, mapping speed, or conversely, integration time to reach a given map noise depth, isan important consideration in modern CMB telescope design. The maximum mapping speed ob-tainable for a given telescope design is proportional to the throughput, or number of independentdiffraction-limited single-moded beams that can be simultaneously measured by populating theimage plane with detectors. Throughput is quantified by the effective area A multiplied by thesolid angle Ω of the combined beams at a given plane in the system, and is a conserved quantityat any position along the optical path. Mapping speed is also proportional to the optical efficiencyof the system, which is the fraction of power in a given spatial mode on the sky that is transmitted

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to the detector. Optical efficiency is negatively impacted by absorption or scattering in the opticalelements, including mirrors, cryostat window materials, lenses, feeds, and optical structures onthe detectors themselves. Mapping speed is also reduced by thermal noise introduced by emissiveoptical elements and other surfaces in the telescope, hence it is desirable to cool emissive elementssuch as refractive optics and stops.

2.1 Small aperture telescope designsSmall CMB telescopes, with apertures < 1 m, are designed to probe the inflationary gravitationalwave B-mode signal at large angular scales, ` < 200, but not the delensing signal at l ¿ 200. Smalltelescopes reviewed here include ABS, BICEP3, CLASS, EBEX, Keck Array/Spider, and Piper.BICEP3 and Keck Array/Spider use all refractive elements, whereas the other experiments usedual off-axis reflector designs, and (with the exception of ABS) cold refractive reimaging optics.

BICEP3 is an evolution of the Keck Array/Spider optical design. Both are a simple two-lens design with a stop. The lenses and stop are at 4 K. The BICEP3 telescope and each KeckArray/Spider are single-frequency, which simplifies the anti-reflection (AR) coating implementa-tion. Multi-frequency coverage is accomplished via the use of multiple telescopes. Lenses in BI-CEP2/Keck Array are fabricated from HDPE plastic, whereas those in the larger 520 mm apertureBICEP3 telescope are alumina, which is both higher index and lower loss than HDPE, resulting inthinner more transmissive lenses.

Discussion of CLASS/PIPER here.Discussion of EBEX.Discussion of ABS.

2.2 Large aperture telescope designsDiscussion of Advanced ACTPol, EBEX, QUIET, Simons Array, SPT-3G here.

2.3 Table of current telescope/instrument designsPossible items for table: aperture, f/#, min Strehl ratio, f-lambda at 150 GHz, A*Omega, ?

See Google Docs page linked here for current table status.

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3 Design studies for next generation telescopesAs described in the CMB-S4 science book, several of the science goals require arcminute-scaleresolution, which roughly translates to telescope apertures between a few and ten meters. Thisrequirement has motivated the study of new optics design concepts. Designs with lower levels ofsystematics (e.g., cross polarization) and larger throughput than existing telescopes are of particularinterest to illuminate much larger numbers of detectors than current observatories.

In addition to these “large-aperture” telescopes, CMB-S4 may also require “small aperture”(≤ 1 m) telescopes to characterize larger angular scales. Due to the lower costs associated withbuilding and operating smaller telescopes, there is greater incentive to develop new designs for thelarger telescopes, which is why we focus on larger telescopes here.

One optics design concept has been shown to achieve substantially larger optical throughputthan existing off-axis Gregorian telescope designs, combined with reduced systematics [3]. Thisconcept is based on a crossed-Dragone design with higher f/# than had been studied previously.Specifically, previous studies of crossed-Dragone CMB telescope designs focus on telescopes withfocal ratios closer to f/1.5 with detector arrays at the telescope focus [8, 9, 10, 11]. For ground-based telescopes controlling spillover with this approach generally requires either severely under-illuminating the primary mirror [9] or cooling the entire telescope to cryogenic temperatures [8],which is not practical for a large telescope.

Figure 1: Left: Optical throughput comparison for crossed-Dragone telescopes with different f/#and aperture and the off-axis Gregorian ACT design [3]. Right: Preliminary high-throughput CDtelescope designs [3]. The CCAT-prime telescope design (Fig. 2) evolved from these.

The new crossed-Dragone concept instead controls spillover past the mirrors using cryogenicrefractive optics, similar to those used in existing large-aperture telescopes [12, 13, 14]. Refractiveoptics naturally couple to telescopes with larger focal ratios, which also increases the availableoptical throughput as shown in Fig. 1. Two six-meter telescope designs following this approach

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are also shown in Fig. 1. These telescope designs are capable of providing a diffraction-limitedfield of view for roughly 10× more detectors than will be deployed on upcoming “Stage III”telescopes [3].

3.1 Systematic considerations and the CCAT-prime concept

Figure 2: A preliminary design for the CCAT-prime telescope provides a 180 elevation rangewithout tilting the cryogenic instruments. This is accomplished by rotating the telescope in eleva-tion along the optical axis between the secondary and flat tertiary, and is shown at 45 and 135

elevation. The extended elevation range enables the equivalent of one telescope bore sight rotationat each observing elevation. These two telescope bore sight positions could be combined with aninstrument rotator that provides arbitrary bore sight rotations for the cryogenic instrument.

The crossed-Dragone telescope design described above has recently been adopted by the CCAT-prime project (www.ccatobservatory.org) and is being studied in greater detail as a candidate tele-scope design for CMB-S4. A preliminary CCAT-prime engineering study is shown in Fig. 2. Thecompact nature of this design enables an unusual optical layout in which the elevation axis isaligned with the optical axis between the secondary and tertiary, providing a single Nasmyth po-sition for instruments. A fixed, flat tertiary folds the focal plane down, keeping the overall sizecompact while improving stability for heavy instruments by shifting them down, closer to the az-imuth platform. The instruments only rotate with the azimuth structure; they do not tip in elevation.This simplifies instrument design and could allow for a simple instrument rotator to help with sys-tematics. Due to the symmetric nature of the telescope mount, the bore sight can be flipped on thesky by rotating in elevation beyond zenith (> 90 deg elevation, Fig. 2) and coming back around180 deg in azimuth. Some baffling is inherent in the structure, as the optics are mounted insideit, and more baffling or a co-moving ground screen would be straightforward to add. The designalso lends itself to having an integrated shutter to provide protection from weather during poorobserving conditions.

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4 Telescope engineering to improve systematicsCMB-S4 will require exquisite control of systematic errors, so the telescopes must be designed tohave low pickup, stable optics, and good pointing while scanning fast enough to freeze atmosphericbrightness fluctuations. It may also be necessary to measure systematic errors that are fixed relativeto the instrument, e.g., by rotating the camera or the entire telescope about boresight. CMB-S4will build on experience with existing telescope designs, but the scale of CMB-S4 may allowapproaches that were deemed impractical for current experiments. Some of these approaches aredescribed below; all will require design and manufacturing studies to assess their viability forCMB-S4.

Exactly what can be tolerated for pickup outside the main beam requiress some investigation,but to give a sense of what will be needed, the EPIC 1.4m design has ∼ −80 dB, ∼ −20 dBisidelobes, and gives ∼ 0.1 nK rms polarized pickup from the galaxy at 150 GHz (“Study of theExperimental Probe of Inflationary Cosmology - Intermediate Mission for NASA’s Einstein Infla-tion Probe,” NASA, 4 June 2009). There are two approaches for controlling pickup: (i) reducescattering, which means low blockage in the obvious sense of off-axis optics, enough clearanceto avoid sidelobes due to clipping the beam, and smooth optical surfaces to avoid scattering fromgaps between mirror segments; and (ii) control what does get scattered, which requires reflectingshields and/or absorbing baffles to eliminate pickup in far sidelobes.

The pointing requirements for CMB-S4 will be stringent: something like beamwidth/100 or1.5” (W. Hu, M. M. Hedman and M. Zaldarriaga, “Benchmark parameters for CMB polarizationexperiments,” Phs. Rev. D 67, 043004 (2003)), so the telescope structures must be stiff. Stableoptics also require stiff structures, and schemes to keep the optical surfaces free of water, snow,and ice.

4.1 Monolithic mirrorsFabrication of a monolithic, millimeter-wavelength mirror larger than a few meters in diameteris challenging, so all existing, large, CMB telescopes (i.e., ACT and SPT) have mirrors made of∼ 1 m segments with ∼ 1 mm gaps between the segments. Scattering from the gaps generatessidelobes which account for ∼ 1% of the telescope response. It is difficult to make the gapssmaller because some clearance is needed for assembly and manufacturing tolerances. Variousgap cover/filler schemes have been tried, but a robust solution has not been demonstrated.

Monolithic mirrors were not practical for the large, stage-3, CMB telescopes, but CMB-S4 willinvolve multiple telescopes, so the cost of developing a fabrication approach for large monolithicmirrors may be reasonable. There is obviously no point pursuing monolithic mirrors unless therest of the telescope design is consistent with small sidelobes.

The key issues for monolithic mirrors are: (i) fabrication errors; and (ii) thermal deformation.Figure 3 shows surface error contributions for a monolithic, aluminum mirror, which is an

obvious choice for low cost. A 5 m diameter, λ = 1 mm mirror seems possible if thermal gradi-ents through and across the mirror are < 1 K, which is what the ∼ 1 m diameter × 50 mm thickCSO primary mirror segments achieve at night. Keeping thermal gradients below 1 K in a largealuminum mirror will require insulation on the back of the mirror, a reflective front coating for

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daytime operation, and maybe active control (e.g., cooling the back of the mirror at night). ACFRP mirror would have an order of magnitude better thermal performance, but it might not bepractical to fabricate a large monolithic CFRP mirror with the required surface accuracy.

4.2 Boresight rotationA few experiments (e.g., DASI, CBI, QUAD, QUIET, BICEP, Keck Array) have included boresightrotation to measure systematic errors that are fixed with respect to the instrument (e.g., instrumentalpolarization) and vary slowly (e.g., on timescales of tens of seconds). All these experiments haveor had small telescopes, or arrays of small telescopes; the largest boresight rotator was the 6 mdiameter platform on the CBI. Large telescope projects have generally dismissed boresight rotationas impractical, but it may be needed to achieve CMB-S4 sensitivity levels, and may be reasonablegiven the scale of CMB-S4.

The key issues for boresight rotation are: (i) balancing the telescope structure while also pro-viding adequate range of motion; and (ii) protecting the drive mechanisms from the weather.

A mount that supports boresight rotation can wrap around the outside of the telescope, whichallows full range of motion with a naturally balanced structure (no counterweight), but resultsin a massive, expensive mount with large mechanisms that are difficult to protect. Alternatively,a compact, inexpensive, enclosed mount can be placed behind the telescope, but this requires acounterweight which results in limited range of motion because the counterweight interferes withthe mount. Figure 4. shows a concept for a compact mount with boresight rotation. The designprovides an optical bench that can support a single, large, off-axis telescope, or an array of smallertelescopes, inside a deep baffle. The compact drive mechanisms can be enclosed, with access frombelow, which is appropriate for a site that has severe snow storms or very low temperatures.

4.3 Shields and bafflesDeep, comoving, reflective shields and/or absorbing baffles will be needed to control pickup in thefar sidelobes. Good shielding is a key reason for the success of small CMB telescopes in makingmeasurements at low `, but a full shield or baffle may also be practical for a large telescope, e.g.,the mount in Figure 4 can accommodate a 5 m telescope inside a deep, cylindrical baffle that issupported by a light, CFRP spaceframe.

The key issues for shields and baffles are: (i) adequate mechanical stability, to avoid time-varying pickup, e.g., due to wind buffeting; (ii) keeping surfaces clear of water, snow, and ice,which change the optical loading; (iii) baffle temperature variations, which cause variations inoptical loading; and (iv) survival of absorbing coatings.

Mechanical stability is more challenging for a reflective shield because any part of the surfacethat sees scattered light must be stable. For an absorbing baffle, the rim must be stable, but the restof the baffle can move relative to the telescope beam as long as the baffle is truly black. There isno practical experience with large absorbing baffles, so the effect of temperature variations needsconsideration. Some work must also be done to identify or develop a light, robust, weather-resistantabsorber.

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1 2 3 4 50

10

20

30

40

50

Mirror diameter (m)

rms

surf

ace

erro

r (µ

m)

gravity

∆T=1K across mirror∆T=1K through mirror

fabrication

10m/s wind

total surface errordiffraction limit at λ=1mm

Figure 3: Surface error vs. diameter for a monolithic aluminum mirror with thickness =diameter/4. Temperature gradients across the mirror change the thickness, while a temperaturegradient through the mirror causes cupping. The gravitational deformation model is the deflec-tion of a simply supported plate, and wind-induced deformation is the gravitational deformationscaled by the ratio of wind pressure to mirror weight per unit area. The fabrication error model has50 µm rms for a 10 m mirror, with error scaling as the square of diameter, combined in quadra-ture with a setup error of 5 µm rms. The model is based on the OVRO 10.4 m primary mirrors,which were machined as a single piece, and the 1 m segments for SPT. The horizontal dashed linecorresponds to 80% Strehl ratio (λ/27 rms surface error) at λ = 1 mm.

Baffle%

Counterweight%

Pier%

Azimuth%bearing%

Boresight%bearing%

Hexapod%

Mirrors%&%

cameras%

Op>cal%bench%

Figure 4: Concept for a telescope mount with boresight rotation. The mount provides a large,flat optical bench that can accommodate various arrangements of cameras and off-axis mirrors.Standard, slewing-ring bearings allow fast scanning in azimuth and rotation about boresight tomodulate/measure polarization. Zenith angle motion is controlled by a hexapod that provides a stiffconnection between the azimuth and boresight slewing rings. The blue structure is a lightweightCFRP spaceframe that supports an absorbing baffle to reduce pickup. Dimensions in the diagramare based on a 6 m diameter optical bench and 2 m diameter slewing rings.

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5 Potential future studies and development areasThe S4 effort would profit from a variety of studies with respect to telescope and receiver design.A comprehensive list is beyond the scope of this document, but here are a few categories and itemsto be considered in the course of decisions about and designs of the S4 telescopes.

Size of Telescopes and Receiver components

1. incremental cost of multiple telescopes of same design versus cost of matching telescopeapertures to each frequency

2. for scenarios with both large and small aperture telescopes, is there an optimal resolutionratio between the small and large apertures

3. size of windows and effectiveness of IR shaders

4. efficiency of refractive optics (and resulting changes in instrument sensitivity) versus refrac-tive aperture

5. cost and feasibility of monolithic mirrors versus aperture

6. cost: marginal cost of additional telescopes vs. larger throughput, including mount costs

7. cost and efficiency: multiple vs. single telescopes per mount

8. cost and efficiency: multi-frequency vs. single band optics

Number of Pixels and Receivers per Telescope

1. efficiency losses caused by downtime from malfunctions or upgrades

2. the ease of matching frequencies to the size of telescope

3. net design effort to balance throughput and such things as sidelobe mitigation

4. marginal weight increase due to receivers as it affects total weight and therefore mount cost

5. nonlinearity of receiver cost as a function of pixel count

Large Telescopes — Crossed Dragone vs. Gregorian (or Cassegrain)

1. sidelobe mitigation

2. spillover mitigation

3. room for receiver(s)

4. mechanical stiffness, effort needed for flexure mitigation, possibility for boresight rotation

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5. accomodation for waveplate(s)

6. throughput

7. new telescope concepts

Small Telescopes — Refractor vs. Crossed Dragone (vs. Greg. or Cass.)

1. sidelobe mitigation

2. spillover mitigation (especially: all cryogenic or oversized mirrors for reflector)

3. overall complexity – alignment, stability, etc.

4. integration into mount with boresight rotation

5. cost: mount size/complexity with boresight rotation

6. cost and efficiency: antireflection requirements

7. throughput

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A Optics designs for current projects

A.1 Advanced ACTPol

Receiver Cabin

Secondary structure

Receiver

Elevation frame

Receiver support structure

Primary structure

2 m Array

LPE

Lens 3

2X LPE Lyot stop

Lens 2 Lens 1

300K filters 40K filters

4K filters

Figure 5: Left: ACT telescope optics and mechanical structure. Right: Raytrace of AdvancedACTPol receiver optics, which includes three optics tubes: one on top and two symmetric tubes onbottom [15].

Advanced ACTPol (AdvACT) is the third instrument upgrade for the 6 m Atacama CosmologyTelescope (ACT). The 6 m primary and 2 m secondary are arranged in a compact off-axis Gregorianconfiguration to give an unobstructed image of the sky. The details of the telescope optics designare presented in [4], while the ACTPol and Advanced ACTPol receiver optics designs are presentedin [15, 16]. Figure 5 shows a raytrace through the ACT mechancial structure as well as through theAdvanced ACTPol receiver optics. Illumination of the primary mirror is controlled using a 1K Lyotstop. To minimize ground pick-up during scanning, the telescope has two ground screens. A large,stationary outer ground screen surrounds the telescope and a second, inner ground screen connectsthe open sides of the primary mirror to the secondary mirror, and moves with the telescope duringscanning.

ACTPol and Advanced ACTPol use the same receiver with three independent optics tubes.Both use large silicon lenses with two and three layer metamaterial antireflection (AR) coatingsfor silicon lenses [17]. These coatings offer the advantages of negligible dielectric losses (< 0.1%),sub-percent reflections, polarization symmetry equivalent to isotropic dielectric layers, and a per-fect match of the coefficient of thermal expansion between coating and lens. Each optics tubefocuses light onto a two-frequency multichroic detector array at one of the following frequencypairs: 28&41 GHz, 90&150 GHz, or 150&230 GHz [13]. The AdvACT reimaging optics havef/1.35 at the array focus. A pixel-to-pixel spacing of 4.75 mm in the recently deployed 150&230GHz array leads to approximately 1.8&2.5 f-λ spacing. A UHMWPE vacuum window is usedcombined with metal mesh filters to control out of band radiation.

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A.2 BICEP3

Figure 6: Left: BICEP2/Keck Array ray trace Right: Bicep3 raytrace.

BICEP3 is a cryogenic refractor of aperture 0.52m, with 2 alumina lenses ([18, 19]) in an f/1.6system. The main cryostat volume is 29in in diameter x 95in high. It operates at 95 GHz on a3-axis mount at the South Pole. The field used is 14.1 half-opening angle, nearly the unvignettedfield of view. At 2mm wavelength the design gives Strehl > 0.96 over the full unvignetted field(top-hat illumination assumed).

The lenses are 99.8% pure alumina. The lenses and stop are at 4K. The window is HDPE,filters consist of metal mesh filters at (nom.) ambient, alumina filter at 50K, another mesh filterbelow that, 2 Nylon filters at 4K, and Ade edge filters at 250mK over each detector module. Thewindow, alumina components, and Nylon filters are single-layer AR coated; the alumina AR isan epoxy mix. A co-moving absorptive forebaffle and a reflective groundshield mitigate groundsource contamination.

2016 95 GHz light detector count is 2400. Pixels are Bock phased slot antenna arrays withtapered weighting to approximate Gaussian beams ([20]). The 1/f noise knee after atmosphericcommon-mode rejection from detector pair differencing is well below the degree-scale scienceband ([21, 2]). Beam systematics are averaged down by boresight rotation and residual temper-ature to polarization beam leakage is removed by deprojection ([22]). Thus, a (fast) polarizationmodulator is not used in BICEP3 (as with BICEP2 and Keck).

The mount (origially built for BICEP1) provides EL down to ∼ 50, full AZ, and boresight

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rotation 255. The latter provides for two 45 offset pairs of 180 complement boresight angles (4angles total) for full Q/U discrimination and cancellation of several beam related systematic errors.Mapping is performed with a sequence of constant EL scans at 2.8/s in AZ.

The BICEP-Array receivers will be substantially the same as BICEP3.

A.3 CLASS

UHMWPE Window

Figure 7: Left: CLASS system overview. Right: CLASS site rendering, showing the two mountswith four telescopes.

The Cosmology Large Angular Scale Surveyor (CLASS) consists of four telescopes sharingsimilar optical layouts [23]. One telescope operates at 40 GHz, two at 90 GHz, and the final tele-scope is a dichroic 150/220 GHz, hereafter the high-frequency (HF) telescope. A 60-cm-diametervariable-delay polarization modulator (VPM) is the first element in the optical chain, providing ∼10 Hz front-end polarization modulation [24]. Ambient-temperature, off-axis, elliptical, 1-meterprimary and secondary reflectors reimage the cold stop of the receiver at 4 K onto the VPM. Cryo-genic reimaging lenses, one at 4 K and one at 1 K, focus light onto the focal plane of feedhorn-coupled transition-edge-sensor bolometers. The CLASS design emphasizes per-detector efficiencyand sensitivity with 10 dB edge-taper illumination of the cold stop. The CLASS telescopes providediffraction-limited performance over a large, 20 field of view with resolutions ranging from 90′

at 40 GHz to 18′ at 220 GHz. Three-axis mounts give azimuth, elevation, and boresight rotations,with two telescopes on each of two mounts (See Figure 7). Co-moving ground shields and bafflesreduce ground pickup.

The lenses for the 40 and 90 GHz telescopes are made of high-density polyethylene (HDPE),while the HF telescope employs silicon lenses. All of the lenses are anti-reflection (AR) coated

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with simulated dielectrics cut directly into the lens material. The receivers have vacuum windowsapproximately 50 cm in diameter made of ultra-high molecular-weight polyethylene (UHMWPE).A combination of capacitive-grid metal-mesh filters and absorptive PTFE and nylon filters rejectinfrared radiation.

A.4 EBEX

Figure 8: EBEX optical design raytrace schematic consisting of two ambient temperature reflec-tors in an off-axis Gregorian configuration and a cryogenic receiver (left). Inside the the receiver(right), cryogenically cooled polyethylene lenses formed a cold stop and provided diffraction lim-ited performance over a flat, telecentric, 6.6 field of view. A continuously rotating achromatichalf-wave plate placed near the aperture stop and a polarizing grid provided the polarimetry capa-bilities.

The EBEX telescope was a balloon-borne CMB polarimeter observing at 150, 250, and 410 GHz.It was required to have flat telecentric focal planes, a large diffraction limited field of view definedas Strehl ratio > 0.9, a cold stop to control sidelobe response, as well as a continuously rotatingachromatic half-wave plate [25] and polarizing grid to provide polarimetry, all while remainingsufficiently compact to fit on a balloon payload [26].

To achieve this the EBEX optical system consisted of a 1.05 m, f/1.9, ambient temperature,Gregorian Mizuguchi-Dragone [6, 7] reflecting telescope and a cryogenic receiver containing 5ultra-high molecular weight polyethylene re-imaging lenses, see Figure 8. The mirrors were over-sized to suppress sidelobe pick-up; the illuminated aperture is 1.05 m while the physical apertureis 1.5 m. The reimaging lenses preserved the f-number of the system while forming a 1 K coldstop, the location of the continuously rotating achromatic half-wave plate, enlarging the diffractionlimited field of view to 6.6, and forming two flat, telecentric focal planes [26, 27]. On the focalplanes conical feedhorns coupled the detectors, TES bolometers, to free space. Each focal plane

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consisted of 7 wafers, 4 at 150 GHz, 2 at 250 GHz, and 1 at 410 GHz. Each wafer contained 128usable detectors; the system was readout limited [28]. Observing bands were defined by reflec-tive filters above the feedhorns and cylindrical waveguides between the feedhorns and bolometers.Reflective IR filters and one absorbtive Teflon filter were used to reduce load on the cyrostat [29].

A.5 Keck/SpiderKeck and SPIDER are close relatives of BICEP2. Both consist of multiple cryogenic refractors ofaperture ∼ 250mm and f/2.2 of essentially the same optical design as BICEP2 ([30]). Both use JPLdual-polarization slot antenna array coupled TES bolometers ([20]).

Keck consists of 5 telescopes co-aligned in their ground-based mount at the South Pole, eachin its own independent vacuum jacket. Individual telescopes have been assigned each observingseason to different frequency bands from 95 to 230 GHz ([31]). Apertures are 264mm and fieldsof view 15 ([32]).

The Keck telescopes have 12cm thick Zotefoam windows, 50K PTFE and Nylon filters, 4KHDPE lenses and a Nylon filter, and Ade edge filters ([31, 32]). The lenses and filters (exceptthe edge filter) are single-layer AR coated, matched to the frequency band of the detector in use.The stop is at 4K, on the bottom of the first lens. Absorptive comoving forebaffles surround eachtelescope aperture, and along with a reflective groundscreen minimize ground pickup.

The Keck array is on a 3-axis mount (built for DASI). Mapping is performed by a sequence ofconstant EL scans at each of 8 boresight rotation angles, 4 pairs of 180 complements for completeQ/U discrimination and mitigation of beam systematics. AZ scan speed is 2.8 /s. The 1/f noiseknee after atmospheric common-mode rejection from detector pair differencing is well below thedegree-scale science band ([21, 2]). Beam systematics are averaged down by boresight rotationand residual temperature to polarization beam leakage is removed by deprojection ([22]). Thus, a(fast) polarization modulator is not used in Keck (as with BICEP2 and BICEP3).

SPIDER is a balloon experiment with 6 co-aligned telescopes in one large vacuum jacket ([33]).It has some optical differences from Keck (and BICEP2) due to the lower sky loading at altitude(as well as mechanical and thermal differences due to weight and cooling system requirements).

HERE some specific differences – predominantl reflective (vs. absorptive) filter stack for lowcryogenic loading, 1.6K sleeve and stop, lower-G detectors, very aggressive magnetic shielding

The SPIDER telescope have apertures of 250mm and 20 fields of view. The window is 1/8inUHMWPE. Metal mesh filters at ambient, 120K and 30K are above a 4K sapphire half wave plate,Nylon filter, and HDPE lenses ([?]). Between the lenses is a 1.5K cooled sleeve with optical baffles.Between the 2nd lens and the 300mK focal plane is a low pass filter at 1.5K. The dielectrics aresingle-layer AR-coated.

The detectors (for the first circum-Polar flight in January, 2015) are JPL slot antenna array-coupled TES bolometers, with 95 and 150 GHz bandpasses. The detector 1/f noise knee is low([33]), the science goals for 10 < l < 300 are accomodated with available scan rates (see below),and fast polarization modulation is not needed, as with Keck and the BICEP2 and 3 instruments(see also [34]).

The second circum-Polar flight is planned to include NIST feedhorn arrays and OMT-coupleddetectors ([35].

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The gondola provides for AZ and EL scanning. It does not have boresight rotation, so dependson the cryogenic waveplates for Q/U discrimination and slow polarization modulation, with acontribution as well from sky rotation. The scan strategy ([33, 36, 37]) is quite different from Keckand the BICEP telescopes due to gondola inertial limitations; a sindusoidal scan was adopted inboth AZ and EL with variable amplitudes. The waveplates are stepped 22.5 every 12 siderealhours. This generates good mapping coverage, Q/U descrimination, and cross-linking.

A.6 Piper

Figure 9: PIPER. Left: Raytrace of original PIPER optics design. Right: Current PIPER imple-mentation.

PIPER is a balloon-borne instrument to observe CMB polarization at 200, 270, 350 and 600 GHz[38]. Twin co-pointed telescopes survey Stokes Q and U . Like CLASS, the first optical element ofeach telescope is a variable-delay polarization modulator (VPM). The VPM separates sky signalfrom instrument drifts by modulating the incoming polarized signal at 3 Hz, aiding reconstructionof the polarized CMB sky on largest angular scales. The VPM efficiently mitigates instrumentpolarization systematics by being the first optical element. Each of the PIPER VPMs have a 40 cmclear aperture with 36µm wires at 115µm pitch. PIPER uses the bucket dewar from ARCADE,which carries 3000 L of liquid helium. Helium boiloff allows operation without emissive win-dows. Superfluid fountain effect pumps draw LHe to cool all optics to 1.4 K. Cold optics and thelack of windows reduce photon noise and allow PIPER to take full advantage of the float condi-tions (especially at high frequencies) and to conduct logistically simpler, conventional flights fromPalestine/Ft. Sumner and Alice Springs. Each flight is optimized for one band, and flights fromthe northern and southern hemispheres cover 85% of the sky.

Two aluminum mirrors image a 12 cm diameter cold aperture stop (1.4 K) onto the centralregion of the front-end VPM. The entrance pupil is 29 cm in diameter and is undersized to limitedge illumination of the VPM (33 dB edge taper). The stop is a corrugated stack of Eccosorb. The1.4 K environment of the bucket dewar mitigates stray light and acts as a comoving ground screen.

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The reflective fore-optics feed silicon re-imaging optics that use metamaterial anti-reflection layers[17]. The off-axis nature of the fore-optics creates aberrations that can be corrected by de-centeringthe reimaging lenses. The reimaging lenses remain planar to the stop and are oversized to retaincylindrical symmetry for diamond turning. The final lens focuses light onto a 32 × 40 free-spacebackshort-under-grid detector array at f/1.6. The resolution at 200 GHz is 21′ and its Airy diskspans approximately six bolometers.The minimum Strehl ratio within the 6 × 4.7 degree FOV is0.97 [39]. PIPER uses a common detector array for all frequencies. Between flights, the VPMthrow, band-defining filters, and (when necessary) lenses are swapped. This strategy is facilitatedby a backshort that is optimized for 200 GHz and is less efficient at high frequencies where theatmosphere and dust emission are brighter. A narrower passband toward higher frequency alsolimits loading.

A.7 Simons Array

Figure 10: Ray tracing diragram of the POLARBEAR-2 and the Simons array optics. Secondarymirror and cryogenic receiver are shown. Length of cryogenic receiver is 2-meters. Diameter ofthree cryogenic lenses are 500 mm.

The three telescopes that comprise the Simons Array are identical o-axis Gregorian designsthat utilize a 2.5m monolithic primary mirror [14]. The telescope and receiver optics are designedto provide a at, telecentric focal plane over a wide diraction-limited eld of view. The angularresolution of the telescope is 5.2, 3.5, and 2.7 at 95, 150 and 220 GHz respectively. Relativepositions of the primary mirror and the secondary mirror obey Mizuguchi-Dragone condition tominimize instrumental cross-polarization [40]. Each telescope has a co-moving shield to prevent

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side lobe pickup from ground emission and an optical bae around prime focus to block stray lightfrom reaching the window and scattering into the receiver. The rst telescope comprising the SimonsArray - the Huan Tran Telescope (HTT) - was installed in Chile in 2011 and has been operatingnearly continuously with the POLARBEAR-1 experiment since. The second and third telescopeswere installed in early 2016.

The receivers have windows made out of laminated 10-inch thick Zotefoam. Radio-TransmissiveMulti-Layer Insulation (RT-MLI) and 2-mm thick anti-reflection coated alumina plate are anchoredto 50-Kelvin stage as infrared filter [41, 42]. First POLARBEAR-2 receiver, the POLARBEAR-2a, will deploy with the ambient temperature continuously rotating half-wave plate (HWP) [43].Second and third POLARBEAR-2 receivers, the POLARBEAR-2b and the POLAEBEAR-2c, willhave cryogenically cooled HWP at 50-Kelvin stage. All three receivers have three 500 mm diame-ter alumina re-imaging lenses cooled to 4-Kelvin. High index of refraction of alumina allowed foroptics design with lenses that have moderate radius of curvature. The first lens is double convexlens. The second lens and the third lens are plano-convex lenses. Optics design has a cold stopbetween the second lens (aperture lens) and the third lens (collimator lens). Meta-material infraredblocking filters and Lyot stop are mounted at the stop. Final F/# of at the focal plane is 1.9. Op-tics design proves diffraction limited illumination that extends over 365 mm diameter of the focalplanes. Strehl ratio at the edge of focal plane is 0.95 for 95 GHz and 0.85 for 150 GHz.

A.8 SPT-3GSPT-3G is a third generation wide-field trichroic (95, 150, 220 GHz) pixel camera for the SouthPole Telescope (SPT), a 10-m off-axis Gregorian telescope first fielded in 2007 [12, 5]. The newoptical design features a 4 K Lyot stop with a 7.5 m primary illumination, Strehl ratios > 0.97 at220 GHz across the 1.9 degree linear field of view (FOV), and angular resolutions of 1.4′, 1.0′, and0.8′ at 95, 150, and 220 GHz, respectively. The telescope is fitted with a reflective primary guardring, side shields, and a prime focus baffle to mitigate far sidelobe pickup.

SPT-3G employs a new optical design consisting of a warm 1.8-m off-axis ellipsoidal sec-ondary mirror positioned at the Mizuguchi-Dragone angle [6, 7], tertiary folding flat mirror, and asingle receiver with three 720-mm diameter 4 K plano-convex alumina lenses which serve to forma 4 K Lyot stop and a flat telecentric f/1.7 focal plane, see Figure 11. The usable FOV is limitedby vignetting from the lens apertures, not optical aberrations. The 6.8-mm pixel pitch translatesinto pixel spacing of 1.3fλ, 2.0fλ and 2.9fλ at 95, 150, 220 GHz, respectively, and was chosento optimize mapping speed given readout constraints.

The vacuum window is 600-mm inner diameter high density polyethylene (HDPE) with atriangular-grooved anti-reflection (AR) coating. IR filters consist of multiple layers of closed cellpolyethylene foam behind the window, an alumina IR blocker and metal mesh IR shaders at 50 K,and low-pass metal mesh IR filters at 4 K and 300 mK. The lenses are three-layer AR coated usingalumina plasma spray [44] and laminated expanded PTFE.

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SPT-3G Strehl Field Map, 220 GHz

Image Plane X-Position (deg) Im

age

Plan

e Y-

Posit

ion

(deg

)

SPT-3G Optics Ray Trace Diagram

Ellipsoidal Secondary

Flat Tertiary

HDPE Window

Alumina Lenses Lyot Stop

Image Plane

Alumina IR Blocker

Prime Focus

1 m

Figure 11: Left: SPT-3G optics ray trace diagram. The new optics include a warm secondary andflat tertiary mirror, and three cold 4 K alumina lenses which serve to form a 4 K Lyot stop and aflat telecentric f/1.7 focal plane. Right: Strehl ratio field map of the 1.9 degree diameter imageplane at 220 GHz. The Strehl ratio is > 0.97 at 220 GHz across the usable 430-mm image planediameter.

B Recent projects using crossed-Dragone telescopes

B.1 ABSThe Atacama B-Mode Search (ABS) telescope consists of 60-cm, cryogenic primary and sec-ondary reflectors in a crossed-Dragone configuration held at the 4 K stage of the receiver (Fig-ure 12) [8]. This optics design was chosen for its compactness for a given focal-plane area andlow cross-polarization. The reflectors were machined out of single pieces of aluminum. A 25-cmstop at 4 K limits illumination of warm elements. The reflectors couple the 25-cm-diameter ar-ray of 240 feedhorn-coupled, polarization-sensitive, transition-edge-sensor bolometer pairs (480detectors) operating at 145 GHz to the sky with 33′ FWHM beams over a 20 field of view. Thetelescope is an f/2.5 system. The polarization directions of the detectors within groups of ten adja-cent detectors were oriented to minimize cross-polarization and each group was tilted to minimizetruncation on the cold stop. Although neither the orientation of the ten elements within a group northe orientation of different groups was parallel, the detectors are largely sensitive to polarizations±45 to the symmetry plane of the optics.

An ambient-temperature, 33-cm-diameter continuously-rotating half-wave plate (HWP) is placedat the entrance aperture of the receiver [45]. The HWP is made of 3.15-mm-thick α-cut sapphireanti-reflection (AR) coated with 305 µm of Rogers RT-Duroid 6002, a fluoropolymer composite.

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A A

Half-Wave Plate

Vacuum Window

SecondaryMirror

4K

FocalPlane

Filter Stack

PrimaryMirror

40K

G10 (Thermal Isolation)

PulseTube

HeliumFridge

ColdStop

4K Cryoperm Shell

Mu-MetalShield

Baffle

Figure 12: Left: Ray trace of ABS optics. Right: Overview of the ABS Receiver

An air-bearing system provided smooth rotation of the HWP at 2.55 Hz and polarization modu-lation in the detector timestreams at 10.20 Hz. Infrared blocking is provided by capacitive-gridmetal-mesh filters patterned on 6-µm mylar with grid spacings of 150 and 260 µm, along with ab-sorptive 2.5-cm PTFE filters AR coated with porous PTFE at 4 K and 60 K. A 0.95 cm nylon filterAR coated with porous PTFE at 4 K provides additional filtering below 1 THz. The receiver hasa 3-mm thick ultra-high molecular weight polyethylene (UHMWPE) vacuum window AR coatedwith porous PTFE. A reflective baffle, shown in Figure 12, and a co-moving ground shield reduceground pickup.

B.2 QUIETQUIET was a crossed Dragone telescope and receiver sited in Chile with 1.4m mirrors [9, 46]. Itoperated with 42 and 90 GHz receivers using corrugated feedhorns (19 and 91 feeds, resp.) andno tertiary optics. It did not have a stop above the primary mirror; the absorptive entrance aperturewas large enough to miss any ray-traced beam from the receiver and only intercepted scattered orstrongly diffracted radiation. Above the entrance aperture an absorptive fore-baffle caught severalknown sidelobes.

At a wavelength 2.0mm, the design would give a Strehl > 0.8 field size (assuming uniformillumination at f/1.65) of ∼ 6.6 half-angle. When considering realistic detector beams the Strehl-limited field size is considerably larger.

The receivers had slow feeds to minimize spillover through the telescope, with FWHM of 7 −8.6 (for the two edges of the 90 GHz band, similar for the 42 GHz band). The resulting beam sizes

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Figure 13: QUIET raytrace.

on the sky were 0.5 for the 42 GHz band and 0.22 for the 90 GHz band. The unvignetted f-ratiofor the 90 GHz receiver’s feed locations was 1.65 (full angle 33.7), resulting in less than 0.25%spillover for any feed in the 91 pixel 95 GHz receiver (modeled, not measured). The telescope wassurrounded by a box of Eccosorb (HR-10 on sheet aluminum, protected by Volara foam), so allspillover was intercepted at ambient temperature except the small percentage that made it throughthe entrance baffle onto the sky or back into the receiver itself.

Cross-polar response of both the telescope and the feed horns was also exceptional [47].A larger receiver with 397 identical feeds in the same hex pattern on an unmodified QUIET

telescope would reach ∼ 2.2% spillover for the edge feeds. However, redistributing the feedpattern and widening the mirrors out of the plane of symmetry would reduce that number. Thedesign is not Strehl-limited.

The QUIET telescope was operated on the CBI mount, with 3 axes including boresight rotation.

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