CMB S-4: Telescope Design Considerations

September 12, 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.

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Contents1 Introduction 1

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

2 Current CMB telescope designs and maturity 32.1 Current small aperture telescope designs . . . . . . . . . . . . . . . . . . . . . . . 42.2 Current large aperture telescope designs . . . . . . . . . . . . . . . . . . . . . . . 52.3 Concept for high throughput large aperture telescope design . . . . . . . . . . . . 6

3 Telescope engineering to improve systematics 93.1 Monolithic mirrors . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 93.2 Boresight rotation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 103.3 Shields and baffles . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 10

4 Potential future studies and development areas 13

A Optics designs for current projects 16A.1 Advanced ACTPol . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 16A.2 BICEP3 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 17A.3 CLASS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 18A.4 EBEX . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 19A.5 Keck/Spider . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 20A.6 Piper . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 21A.7 Simons Array . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 22A.8 SPT-3G . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 23

B Projects using crossed-Dragone telescopes 24B.1 ABS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 24B.2 QUIET . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 25B.3 CCAT-prime . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 26

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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 lessons learned from decades ofmeasurement experience. The ten current-generation experiments and two previous-generationtelescope designs presented in Figure 2 have a wide variety of optical design approaches, includ-ing both refractive and reflective primary aperatures, Gregorian and crossed-Dragone mirror con-figurations, single- and multi-camera systems (also called “tertiary re-imaging optics” or “opticstubes”), and polarization modulation mechanisms including rotating half-wave plates (HWPs), re-flective variable-delay polarization modulators (VPMs), telescope boresight rotation, sky-rotation,or no polarization modulation.

Small aperture designs, . 1 m diameter at mm-wavelengths, are used to probe for the inflation-ary gravitational wave B-mode signal, whereas large aperture designs, & 1 m diameter, are used tomeasure the CMB lensing signal, and & 5 m designs can also measure arcminute-scale secondaryanisotropies.

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 elim-inate the use of feed legs in the optical path, which scatter light into far sidelobes. Sources thatcan introduce scan-synchronous spurious signals include the ground, the sun and moon, and thegalaxy. To reduce these signals, reflective shields and absorbing baffles, both co-moving with thetelescope and fixed to the ground, are used to reflect scattered light to the cold sky, preferably nearthe 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, the cre-ation 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. These includemismatches in gain calibration, beam shape, or bandpass in differenced detectors, and by point-ing errors. Cross-polarization, 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 [25, 26], which eliminates cross-polarizationat the 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, soit is desirable to modulate on the sky side of as much of the optical system as possible.

Mapping speed, or conversely, integration time to reach a given map noise depth, is an impor-tant consideration in modern CMB telescope design. Many factors affect mapping speed, includingoptical throughput (i.e. number of diffraction limited pixels), optical efficiency, thermal noise, andtelescope down time.

The maximum mapping speed obtainable for a given telescope design is proportional to thethroughput, or number of independent diffraction-limited single-moded beams that can be simul-

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taneously measured by populating the image plane with detectors. Throughput is quantified by theproduct of the effective area A and the solid angle Ω of the combined beams at a given location inthe system, and is a conserved quantity at any position along the optical path.

Mapping speed is also proportional to the optical efficiency of the system, which is the fractionof power in a given diffraction limited spatial mode on the sky that is transmitted to the detector.Optical efficiency is adversely affected by many factors, including optical aberrations inherent inthe design (characterized in part by the Strehl ratio), scattering due to fabrication imperfectionsor intrinsic material properties, and undesired absorption or reflection in the optical elements in-cluding mirrors, cryostat window materials, polarization modulators, lenses, feeds, and opticalstructures on the detectors themselves.

Mapping speed is also reduced by thermal noise introduced by emissive optical elements andother surfaces in the telescope, hence it is desirable to cool emissive elements such as refractiveoptics, stops, absorptive baffling, and even mirrors.

Lens materials used in current CMB telescopes include high-density polyethylene (HDPE),ultra-high-molecular-weight polyethylene (UHMWPE), silicon, and alumina. HDPE and UHMWPEhave the advantage of being easy to machine, and the relatively low index of refraction (n ∼ 1.5)compared to silicon (n ∼ 3.3) and alumina (n ∼ 3.1) facilitates anti-reflection (AR) coating,but results in thicker and lossier lenses, particularly for large apertures. Material loss typicallydecreases at low temperatures, adding to the benefit of cooling refractive optical elements. Seethe CMB S4 Broadband Optics white paper for more details on lens materials and AR coatingtechnologies.

Finally, telescope down time can have a significant impact on meeting sensitivity goals. There-fore, mechanical robustness, serviceability, and weatherization are important factors in currentCMB telescope designs.

Small- and large-class telescope designs are discussed separately in the sections below. A tablesummarizing some relevant optical parameters is given in Figure 2.

2.1 Current small aperture telescope designsSmall mm-wavelength CMB telescopes, with apertures . 1 m, are designed to probe the inflation-ary gravitational wave B-mode signal at large angular scales, ` < 200, but not the delensing signalat ` > 200. Small telescopes reviewed here include ABS (0.25 m physical aperture), BICEP3(0.53 m), CLASS (0.6 m), Keck Array/Spider (0.25 m), and Piper (0.39 m). BICEP3 and KeckArray/Spider use all refractive elements, whereas the other experiments use dual off-axis reflectordesigns, and (with the exception of ABS) cold refractive reimaging optics.

BICEP3 (and the future BICEP Array) is an evolution of the Keck Array/Spider optical de-sign. Both are a simple two-lens objective/eyepiece design with a stop. The lenses and stop areat 4 K. The BICEP3 telescope and each Keck Array/Spider telescope are single-frequency, whichsimplifies the anti-reflection (AR) coating implementation. Multi-frequency coverage is accom-plished via the deployment of multiple telescopes. Lenses the Keck Array are fabricated fromHDPE plastic, whereas those in the larger 520 mm aperture BICEP3 telescope are alumina. Anabsortive cylindrical baffle is placed in front of the telescope cryostat, and fixed reflective ground

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shields surround the telescopes. The telescope mount allows for boresight rotation to rotate thepolarization angle sensitivity and check for polarization systematics.

The CLASS ground-based experiment and PIPER balloon-borne experiment have similar opti-cal layouts consisting of a variable-delay polarization modulator (VPM) as the first optical elementat the entrance pupil, two elliptical mirrors, cold refractive reimaging optics, and a cold stop. InCLASS, the VPM and mirrors are at ambient temperature, whereas in PIPER all optics are cooledto 1.4 K by LHe boiloff. CLASS uses HDPE lenses and an UHMWPE window, whereas PIPERuses silicon lenses. Both telescopes have comoving baffles surrounding the mirrors and in front ofthe VPM. Similar to BICEP3 and Keck Array/Spider, Multi-frequency coverage is accomplishedfor CLASS using multiple telescopes (40, 90, and 150/220 GHz telescopes) and for PIPER usingmultiple flights, which simplifies AR coating implementation.

ABS used a compact crossed-Dragone dual mirror design and no tertiary re-imaging optics.The mirrors and focal plane were cooled to 4 K, with a cold stop preceding the primary mirror.An ambient temperature continuously-rotating half-wave plate (HWP) and co-moving baffle wereplaced just above the AR-coated UHMWPE vacuum window.

2.2 Current large aperture telescope designsLarge aperture mm-wavelength CMB telescope designs, & 1 m diameter, can be used to mea-sure the CMB lensing signal at ` > 200, and > 5 m designs can also measure arcminute-scalesecondary anisotropies. Large aperture designs reviewed here include Advanced ACTPol (6 mphysical aperture), EBEX (1.5 m), Simons Array (2.5 m), SPT-3G (10 m), and QUIET (1.4 m).All large telescope designs with the exception of QUIET use ambient temperature dual mirroroff-axis Gregorian configuration and cold refractive reimaging optics to form a cold stop and tele-centric (flat) image plane. In contrast to the small-aperture designs, the large-aperture telescopesare all designed to conduct simultaneous observations in multiple frequency bands.

The Advanced ACTPol (AdvACT) receiver design consists of three independent optics tubesfor each of three frequency pairs (28/41, 90/150, and 150/230 GHz). Each optics tube uses aUHMWPE vacuum window and has three cold AR coated silicon lenses, and a 1 K cold stop tocontrol illumination of the primary. AdvACT will use ambient temperature continously rotatingHWP’s just outside each optics tube’s vacuum window. The ACT telescope has reflective shielding,both co-moving and a stationary ground shield surrounding the telescope.

The Simons Array (three telescopes) and SPT-3G share a similar optical design consisting of asingle receiver (per telescope) with three cold alumina lenses and a 4 K cold stop. SPT-3G uses anHDPE window whereas the Simons Array receivers (called the POLARBEAR-2 receivers) eachuse a 10-inch thick laminated zotefoam window. The first POLARBEAR-2 receiver will have anambient temperature continuously rotating HWP just outside the receiver window, whereas thesecond and third receivers will have 50 K HWP’s inside the cryostat window. SPT-3G does nothave plans to use a HWP. Both the Simons Array telescopes and SPT-3G use a prime focus baffle,and reflective co-moving shields, but no fixed ground shields.

EBEX operated at 150, 250, and 410 GHz using a receiver containing five UHMWPE lenses, aco-located cold stop and continuously rotating HWP at 1 K. Polarization sensitivity was achievedby using a polarizing grid inside the receiver cryostat and two focal planes of non-polarization

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sensitive detectors. Each focal plane consisted of seven detector wafers, with each wafer sensitiveto a single frequency band (150, 250, or 410 GHz), defined by reflective filters above the feedhornsand the cylindrical waveguides between the feedhorns and detectors.

QUIET was a crossed Dragone telescope operating at 42 and 90 GHz. The two mirrors wereat ambient temperature, and no tertiary optics were used in front of the cryogenic focal plane offeedhorns. Unlike ABS, it did not use a stop above the primary mirror. Instead, the feeds weresized to under-illuminate the mirrors to minimize spillover. Sidelobes were also mitigated by usingan absporbing baffle in front of the entrance aperture, and surrounding the telescope.

2.3 Concept for high throughput large aperture telescope designAs 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.

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 [6, 7, 8, 9]. For ground-based telescopes controlling spillover with this approach generally requires either severely under-illuminating the primary mirror [7] or cooling the entire telescope to cryogenic temperatures [6],which is not practical for a large telescope.

Figure 1: Optical throughput comparison for large-aperture crossed-Dragone telescopes with dif-ferent f/# and aperture compared to the off-axis Gregorian ACT design [3].

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The new crossed-Dragone concept instead controls spillover past the mirrors using cryogenicrefractive optics, similar to those used in the existing large-aperture telescopes ACT [11], SimonsArray [12], and SPT [10]. Refractive optics naturally couple to telescopes with larger focal ratios,which also increases the available optical throughput as shown in Fig. 1. This crossed-Dragonedesign has recently been adopted for the CCAT-prime project. Preliminary designs and potentialsystematics advantages of this design are shown in Fig. 14 and discussed in §B.3. When combinedwith close-packed optics tubes, a 6 m telescope based on this design is capable of providing adiffraction-limited field of view for > 105 detectors, which is roughly 10× more detectors thanwill be deployed on upcoming “Stage III” telescopes [3].

It is not yet clear whether other telescope designs are capable of providing a similar scaleof diffraction-limited field of view with similar systematics. For example, adding an optimizedtertiary to more traditional off-axis Gregorian designs has not been studied in detail. As describedin section §4, these are high priority areas to study for CMB-S4.

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Project ABS Keck Array Spider Piper BICEP Array CLASSPhysical aperture (m) 0.25 0.25 0.25 0.39 0.52 0.6Illuminated aperture (m) 0.25 0.25 0.25 0.29 0.52 0.35Telescope f/# 2.5 2.2 2.2 1.55 1.6 2, 2, 1.5, 1.5f/# at detector array (if different) 2.2 2.2 1.6 1.6

Minimum Strehl ratio at 150 GHz 0.96 0.97 (200 GHz) 0.99

f-lambda spacing at 150 GHz 2.6 0.5 2.42A*Omega of illuminated arrays (cm^2 sr) 50 6 92A*Omega with Strehl > 0.8 at 150 GHz 51Field of view per array (deg^2) 315 28 315Number of arrays 1 5 2 (4 supported) 5 4Number of telescopes 1 1 2 5 4

Observation frequencies (GHz) 150 95, 150, 220 200, 270, 350, 600

35, 95, 150, 220/280

38, 93, 147, 218

Detectors on sky per frequency 480 (288, 512, 512) x # arrays per freq

384, 6106,7776, 9408/9408

72, 1036, 1190, 1190

# Frequencies per array ("multichroic-ness") 1 1 1 1,1,1,2 1(40,90), 2(150/220)

Window Material UHMWPE Zotefoam HD-30 UHMWPE None HDPE UHMWPEIlluminated diameter of window (m) 0.28 0.26 n/a 0.68 0.35Lens Material N/A HDPE HDPE Silicon Alumina HDPE, siliconTemperatures of reflective optics (K) 4 1.4 -- 300Temperatures of refractive optics (K) N/A 4 1.4 4 4, 1Temperature of cold stop (K) 4 4 1.4 4 4Temperature of detector arrays (K) 0.3 0.25 0.1 0.25 0.05Year of initial (or partial) deployment 2012 2012 2016 2015 2016Year of full deployment (all frequencies) 2012 2013 2014 2020 2020 2018

Project QUIET EBEX Simons ArrayAdv. ACTPol CCAT-Prime SPT-3GPhysical aperture (m) 1.4 1.5 2.5 6 6 10Illuminated aperture (m) 1.05 2.5 5.6 5.5 7.5Telescope f/# 1.65 1.9 1.9 2.5 3 1.7f/# at detector array (if different) 1.9 1.9 1.35 1.5 1.7

Minimum Strehl ratio at 150 GHz 0.9 0.85 0.8 (1 array), 0.93 (2 arrays) 0.81 0.99

f-lambda spacing at 150 GHz 1.74 1.8 1.8 1.3 2A*Omega of illuminated arrays (cm^2 sr) 180 ~2700 250A*Omega with Strehl > 0.8 at 150 GHz 379 ~3000 370Field of view per array (deg^2) 4 deg on sky 0.8 0.9 1.9Number of arrays 14 1 3 up to 50 1Number of telescopes 1 1 1 1 1

Observation frequencies (GHz) 42, 90 150, 250, 410 90, 150, 220, 280

28, 41, 90, 150, 230 90 GHz - 1 THz 90,

150, 220

Detectors on sky per frequency 7588, 7588, 3794, 3794

88, 88, 1712, 2718, 1006 up to ~10^5 5420,

5420, 5420

# Frequencies per array ("multichroic-ness") 1 1 2 2 2 or 3 3Window Material UHMWPE Zote Foam UHMWPE HDPEIlluminated diameter of window (m) 0.28 0.5 0.31 0.6Lens Material N/A UHMWPE alumina silicon aluminaTemperatures of reflective optics (K) 300 300 300 300 300 300Temperatures of refractive optics (K) N/A 4, 1 4 4, 1 4Temperature of cold stop (K) N/A 1 4 1 4Temperature of detector arrays (K) 0.25 0.25 0.1 0.1 0.25Year of initial (or partial) deployment 2009 (test flight) 2017 2016 2020 2016Year of full deployment (all frequencies) 2013 2017 2018 TBD 2016

Figure 2: Table of telescope and instrument parameters for current projects sorted by primaryaperture. Descriptions of each of these projects are in §A and §B.

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3 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.

3.1 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.

3.2 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.

Another approach under study that offers partial boresight rotation is having the telescopeelevation axis aligned with the chief ray between the secondary and a flat tertiary mirror. Thisapproach offers the additional advantage of not tilting the instruments in elevation and is beingpursued by the CCAT-prime project (see §B.3 for details).

3.3 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 is

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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.

no 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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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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4 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 studies and areas ofresearch that could significantly help in the course of decisions about and designs of S4 telescopes.

Can large aperture ground-based telescopes (> 1 m aperture) effectively measure l ≈ 50?This is critical for considering whether CMB-S4 could be composed of only large aperture tele-scopes, or if it will require a mix of large aperture and small aperture. ACT (§A.1), the SimonsArray (§A.7), and SPT (§A.8) are currently attempting to make measurements at lower l with dif-ferent combinations of half-wave plates and scan strategies, but none have published sufficientlylow-l results to access the full spatial range of CMB-S4 science yet. Some new large-aperturetelescope designs have a field-of-view comparable in size to small aperture telescopes (e.g., ∼7

for CCAT-prime, §B.3), which could help, though this may not be tested in time.

Will any ground-based telescopes be able to access l ≈ 5? CLASS ((§A.3)) and PIPER (aNASA balloon, §A.6) are the only current sub-orbital projects targeting this largest angular scalerange. Measurements of both E-mode and B-mode polarization would be valuable to make onthese large spatial scales. CMB-S4 science can be pursued independently of these measurements,but it may be worth expanding the CMB-S4 science case if these projects produce compellingl ≈ 5 results.

What types of telescopes can provide the most useable throughput per unit cost? It hasbeen shown that crossed-Dragone designs can provide substantially larger throughput than currentlarge aperture telescopes (§2.3). However, current telescopes were not designed with the goal ofmaximizing throughput, and it is not yet known whether modifications to more traditional off-axis Gregorian or Cassegrain designs could achieve similar performance. For example, adding anoptimized tertiary mirror to an off-axis Gregorian might significantly increase throughput. Intro-ducing higher-order correction terms to the mirrors of a conventional crossed-Dragone also has thepotential to increase diffraction-limited throughput.

Are there advantages to designing many different telescope apertures as opposed to only hav-ing one or two telescope apertures? Motivations for this could be to match telescope resolutionat different frequencies and potentially to save cost. The answer for this may differ for “large”versus “small” aperture telescopes; for example, it may make sense to adjust the apertures of tele-scopes< 1 m by frequency, but to only build one or two designs for the more costly larger aperturetelescopes, or vice versa.

What is the cost, feasibility, and importance of having monolithic mirrors on large aperturetelescopes? It is clear that gaps between panels in the primary and secondary mirrors generatesidelobe structure in large aperture telescopes, but it is not yet clear how limiting these sidelobes

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would be for CMB-S4 science. Additional simulations and calculations would help with under-standing this. It is also important to have a realistic assessment of the cost and feasibility of usinglarge aperture monolithic mirrors (§3.1).

Is there a practical limit to cryogenic receiver size? As the physical aperture of receiver win-dows increases, more cooling power is needed to remove the additional radiative loading. Forexample, the receiver design in [3] (Fig. 14) includes 50 optics tubes, and would require a largernumber of pulse tube refrigerators than existing receivers. Improved cryogenic modeling wouldhelp to assess whether building receivers with much larger window areas becomes impractical. Inaddition, it would be valuable to understand how quickly large pulse-tube-cooled receivers couldbe cooled from (and warmed to) room temperature for upgrades and repairs. This could lead to adifferent practical limit to receiver size. On the other hand, an integrated liquid nitrogen pre-coolsystem may be able to significantly mitigate some of these concerns.

What are the tradeoffs between having a single large diameter optics tube per telescope ver-sus having many smaller diameter optics tubes? It is clear that a single large diameter opticstube simplifies the design of the cryogenic receiver (e.g., §A.7, §A.8); however, it limits the band-width of the detectors illuminated by that optics tube and increases loss due to having thicker re-fractive optics. Having many smaller optics tubes enables each one to be optimized for a differentwavelength and relaxes space constraints around each detector array (e.g., §A.1,§B.3); however,it appears to lead to increased cryogenic complexity and requires fabrication of more refractiveoptical elements. Use of many optics tubes also appears to increase the useable field of view [3];however, the SPT-3G refractive optics design provides a substantially larger field of view than theGregorian focus alone (§A.8). It is important to understand the optical design and implementationtradeoffs between these options in more detail.

What are the largest practical vacuum windows, filters, and lenses? As the vacuum windowdiameter increases the thickness (and therefore loss) must also increase. Is there a size where theloss becomes so large that it naturally limits consideration of larger optics tubes?

A related question is: are there window materials with sufficiently low loss that they can alsobe used as room temperature refractive optics? Having a refractive optic at the vacuum windowcan considerably decrease the volume of the cryogenic instrument. This is important for designswith multiple optics tubes, which require close-packed refractive optics to make optimal use of thetelescope field of view. The largest tradeoff in doing this is the additional loss due to the thicknessof the refractive optic. If low-loss refractive optics that can also be used as windows are developed,it will enable more efficient use of the diffraction-limited field of view on large aperture telescopes.

What baffle designs are needed and worthwhile for CMB-S4? Different groups use a rangeof absorptive and reflective baffles. It is important to consider how baffles can be designed moreeffectively, especially for large aperture telescopes (§3.3). This is any area of study that wouldbenefit from significant effort to both control systematics and minimize spillover onto absorptivebaffles, which can also minimizes excess photon noise.

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How important are boresight rotations for large-aperture telescopes? Thus far boresight ro-tations have only been implemented on relatively small aperture telescopes. Additional simulationsincluding a variety of sources of systematics could help address whether the additional cost of im-plementing full boresight rotations or partial boresight rotations (§3.2) on large aperture telescopesis worthwhile.

What can be done to reduce the cost of CMB telescopes for CMB-S4? For example, woulddeveloping a common platform and control system for both large and small aperture telescopes leadto significant cost reductions? Studies that have the potential to reduce costs through developingcommon versatile components that could benefit both large and small aperture telescopes may befruitful.

Does it make sense to use existing or planned telescopes for CMB-S4, or will it be more costeffective to build all new CMB-S4 telescopes? Studies that take into account telescope buildingcosts, operating costs, demonstrated performance, optical throughput, etc. would be valuable forhelping assess the most cost effective approach for developing CMB-S4.

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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 [13].

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 [13, 14]. 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 [15]. 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 [11]. 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 ([16, 17]) 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 ([18]). The 1/f noise knee after atmosphericcommon-mode rejection from detector pair differencing is well below the degree-scale scienceband ([19, 2]). Beam systematics are averaged down by boresight rotation and residual temper-ature to polarization beam leakage is removed by deprojection ([20]). 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 [21]. 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 [22]. 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 [23] and polarizing grid to provide polarimetry, all while remainingsufficiently compact to fit on a balloon payload [24].

To achieve this the EBEX optical system consisted of a 1.05 m, f/1.9, ambient temperature,Gregorian Mizuguchi-Dragone [25, 26] 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 [24, 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 ([18]).

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 ([19, 2]). Beam systematics are averaged down by boresight rotationand residual temperature to polarization beam leakage is removed by deprojection ([20]). 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[15]. 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. The length of the cryogenic receiver is 2 m. The diameterof the three cryogenic lenses are 500 mm.

The three telescopes that comprise the Simons Array are identical off-axis Gregorian designsthat utilize a 2.5m monolithic primary mirror [12]. The telescope and receiver optics are designedto provide a flat, telecentric focal plane over a wide diffraction-limited field 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 the Mizuguchi-Dragone conditionto minimize 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 baffle around prime focus to block straylight from reaching the window and scattering into the receiver. The first telescope comprisingthe Simons Array - the Huan Tran Telescope (HTT) - was installed in Chile in 2011 and has beenoperating nearly continuously with the POLARBEAR-1 experiment since. The second and thirdtelescopes were installed in early 2016.

The receivers have windows made out of laminated 10-inch thick Zotefoam. Radio-TransmissiveMulti-Layer Insulation (RT-MLI) and a 2-mm thick anti-reflection coated alumina plate are an-chored to the 50-Kelvin stage as infrared filter [41, 42]. The first POLARBEAR-2 receiver, thePOLARBEAR-2a, will deploy with the ambient temperature continuously rotating half-wave plate(HWP) [43]. The second and third POLARBEAR-2 receivers, the POLARBEAR-2b and thePOLAEBEAR-2c, will have a cryogenically cooled HWP at 50-Kelvin stage. All three receivershave three 500 mm diameter alumina re-imaging lenses cooled to 4-Kelvin. High index of refrac-tion of alumina allowed for an optics design with lenses that have moderate radius of curvature.The first lens is a double convex lens, whereas the second lens and the third lens are plano-convexlenses. The optics design has a cold stop between the second lens (aperture lens) and the third lens(collimator lens). Meta-material infrared blocking filters and Lyot stop are mounted at the stop.The final F/# of at the focal plane is 1.9. The optics design provides diffraction limited illuminationthat extends over 365 mm diameter of the focal planes. The Strehl ratio at the edge of focal planeis 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 [10, 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 [25, 26], tertiary folding flat mirror,and a single receiver with three 720-mm diameter 4 K plano-convex alumina lenses which serveto form a 4 K Lyot stop and a flat telecentric f/1.7 focal plane, see Figure 11. The usable FOVis limited by vignetting from the lens apertures, not optical aberrations. The 6.8-mm pixel pitchtranslates into pixel spacing of 1.3fλ, 2.0fλ and 2.9fλ at 95, 150, 220 GHz, respectively, and waschosen to 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 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) [6]. 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 [7, 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.

B.3 CCAT-primeThe crossed-Dragone telescope design presented in [3] and described in §2.3 has recently beenadopted by the CCAT-prime project (www.ccatobservatory.org) and is being studied in greaterdetail as a candidate telescope design for CMB-S4. A preliminary CCAT-prime engineering studyis shown in Fig. 14. The compact nature of this design enables an unusual optical layout in whichthe elevation axis is aligned with the optical axis between the secondary and tertiary, providing asingle Nasmyth position for instruments. A fixed, flat tertiary folds the focal plane down, keepingthe overall size compact while improving stability for heavy instruments by shifting them down,closer to the azimuth platform. The instruments only rotate with the azimuth structure; they do nottip in elevation. This simplifies instrument design and could allow for a simple instrument rotatorto help with systematics. Due to the symmetric nature of the telescope mount, the bore sight canbe flipped on the sky by rotating in elevation beyond zenith (> 90 deg elevation, Fig. 14) andcoming back around 180 deg in azimuth. Some baffling is inherent in the structure, as the optics

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Figure 14: Left: Preliminary high-throughput CD telescope design [3]. The CCAT-prime tele-scope design evolved from this. Middle pair: A preliminary design for the CCAT-prime telescopeprovides a 180 elevation range without tilting the cryogenic instruments. This is accomplishedby rotating the telescope in elevation 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 onetelescope bore sight rotation at each observing elevation. These two telescope bore sight positionscould be combined with an instrument rotator that provides arbitrary bore sight rotations for thecryogenic instrument. Right: Concept for a CMB-S4 receiver with 50 optics tubes that could eachilluminate between 2,000-3,000 detectors, providing > 105 detectors on a single telescope.

are mounted inside it, and more baffling or a co-moving ground screen would be straightforwardto add. A clear advantage of having the optics ”buried” inside the mount is lower wind loadingon the mirrors as they are effectively inside an enclosure, and elimination of the typical secondarysupport structure. The design also lends itself to having an integrated shutter to provide protectionfrom weather during poor observing conditions.

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