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FIRST: Fibered Imager foR Single Telescope Feasibility study report ...

FIRST: Fibered Imager foR Single TelescopeFeasibility study report to CFHT© Olivier Lai 2006October 1 st 2008Report prepared by:Jean-Philippe Berger 1 , Pascal Bordé 2 , Olivier Chesneau 3 , Pierre Fédou 4 , Pierre Kervella 4 ,Olivier Lai 5 , Alain Lecavelier 6 , Takayuki Kotani 4 , Sylvestre Lacour 1,4 , Guy Perrin 4 , StephenRidgway 7 , Éric Thiébaut 8 , Alfred Vidal-Madjar 6 , Julien Woillez 91Laboratoire d’Astrophysique de l’Observatoire de Grenoble2Institut d’Astrophysique Spatiale3Laboratoire Fizeau, Observatoire de la Côte d’Azur4LESIA, Observatoire de Paris5Canada-France-Hawaii Telescope Corporation6Institut d’Astrophysique de Paris7National Optical Astronomy Observatory8Centre de Recherches Astronomiques de Lyon9W.M. Keck Observatory


Executive summary................................................................................................................ 41. Introduction..................................................................................................................... 52 Scientific case .................................................................................................................. 62.1 Exoplanets................................................................................................................. 62.2 Debris disks............................................................................................................... 82.3 Young Stellar Objects.............................................................................................. 122.4 Cepheids ................................................................................................................. 142.5 Novae...................................................................................................................... 162.6 Compact nebulae around evolved stars .................................................................... 182.7 From stellar surfaces to mass loss............................................................................ 202.8 Active galactic nuclei .............................................................................................. 253 Science requirements...................................................................................................... 274 Principle of fibered aperture masking ............................................................................. 284.1 The issue of high dynamic range imaging................................................................ 284.2 Aperture masking .................................................................................................... 294.3 Aperture masking with single-mode fibers............................................................... 304.4 Photometric calibration and Image reconstruction.................................................... 314.5 Theoretical performance.......................................................................................... 324.5.1 Sensitivity......................................................................................................... 324.5.2 Dynamic range ................................................................................................. 324.5.3 Field of view..................................................................................................... 335 Instrument concept......................................................................................................... 345.1 Overview................................................................................................................. 345.2 Pupil sampling......................................................................................................... 365.2.1 Sub-pupil sizes ................................................................................................. 365.2.2 Lenslet array..................................................................................................... 375.3 Injection in fibers .................................................................................................... 385.4 Beam combination................................................................................................... 385.5 Fibers ...................................................................................................................... 405.5.1 Visible wavelength range single-mode fibers .................................................... 415.5.2 Input fiber bundle ............................................................................................. 415.5.3 Output fiber assembly....................................................................................... 425.5.4 Dispersion ........................................................................................................ 425.5.5 Polarization ...................................................................................................... 435.5.6 Throughput....................................................................................................... 435.6 Segmented mirror.................................................................................................... 435.7 Differential refraction compensation........................................................................ 445.8 Detection................................................................................................................. 455.8.1 EMCCD ........................................................................................................... 455.8.2 Custom Photon counting detector ..................................................................... 465.8.3 Brief comparison of the two detector types ....................................................... 475.9 Instrument upgrades ................................................................................................ 485.9.1 Spectral resolution ............................................................................................ 485.9.2 Field of view..................................................................................................... 485.9.3 Sensitivity......................................................................................................... 485.9.4 Beam combination............................................................................................ 485.10 Expected performance ........................................................................................... 495.10.1 Summary of instrument and environment characteristics................................. 495.10.2 Sensitivity....................................................................................................... 502


5.10.3 Dynamic range................................................................................................ 515.11 Compliance of baseline instrument with science specifications.............................. 535.12 Laboratory demonstrator........................................................................................ 545.12.1 Overview........................................................................................................ 555.12.2 Injection optics ............................................................................................... 555.12.3 Beam combiner............................................................................................... 565.12.4 Spectrometer and Anamorphic optics.............................................................. 575.12.5 Detector.......................................................................................................... 585.12.6 Preliminary results.......................................................................................... 585.12.7 Conclusion...................................................................................................... 596 Development requirements for the baseline project ........................................................ 616.1 Optical design ......................................................................................................... 616.2 Thermal and Mechanical design .............................................................................. 616.3 Electronics .............................................................................................................. 616.4 Data transfer and storage ......................................................................................... 616.5 Beam combination................................................................................................... 627 Before operations ........................................................................................................... 647.1 Tests in Meudon...................................................................................................... 647.2 Shipping.................................................................................................................. 647.3 Receiving ................................................................................................................ 647.4 Commissioning ....................................................................................................... 648 Operations and data processing ...................................................................................... 668.1 Set-up...................................................................................................................... 668.2 Calibrations............................................................................................................. 668.3 Data acquisition....................................................................................................... 668.4 Data reduction......................................................................................................... 668.5 Observing conditions requirements.......................................................................... 669 Development roadmap and preliminary cost estimate ..................................................... 679.1 Project schedule ...................................................................................................... 679.2 Manpower requirements .......................................................................................... 679.3 Cost estimate........................................................................................................... 6810 References.................................................................................................................... 703


4Executive summaryWe propose an instrument whose purpose is to provide a high angular resolution capabilityfor CFHT at high dynamic range in the visible range of wavelengths. This instrument buildson a concept recently proposed by Perrin et al. (2006): the telescope is turned into a 126 pupilredundant high precision fibered interferometer. The fundamental idea behind this concept isthat phase corrugations are filtered by single-mode fibers as demonstrated by long baselineinterferometers. Redundancy allows to calibrate the residual phase differences and thefluctuations of intensity coupled in each fiber thanks to a self-calibration type imagereconstruction algorithm whose output is both an image of the object and the phase andintensity map of the pupil for each individual exposures. The result is a fix Optical TransferFunction (to within photon noise), the key to high dynamic range imaging. A tremendousadvantage of the technique is that the dynamic range is uniform across the entire imaged fieldand photon noise limited. In particular, very high dynamic ranges are achievable very near acentral bright object, well within the diffraction limit when the object is simple (like amultiple point-like source object). Speckles are totally eliminated (thanks to spatial filteringby single-mode fibers) and images can be stacked to improve the dynamic range (it increaseswith the square root of the total collected flux). This new imager type is undergoing aprototyping phase in Meudon Observatory and first results are presented in this report. Weplan to achieve sky demonstration in Meudon before we start building FIRST for CFHT.Given the high angular resolution potential of CFHT in the visible, 33% better than HST,FIRST provides CFHT with a unique capability in today’s astronomy. The performancesderived for the instrument are outstanding: dynamic range as high as 7.10 6 . A limitingmagnitude of 12.7 can be reached if the instrument benefits from a correction by PUEO. Wetherefore recommend FIRST to be coupled to PUEO. Images are reconstructed in spectralchannels with a resolution as high as 250. Spectral channels can be summed to increasesensitivity and reduce spectral resolution in a still photon-noise limited regime as detectorread-out noise is less than 0.1e - (spectral resolution can be reduced to R=12 to keep maximumfield of view).We show in this report that such an instrument can be built for CFHT at a reasonable cost(682 k$ for our baseline, our preferred concept, 559 k$ for an alternative concept with areduced detector cost but with poorer spectral capability).We demonstrate that FIRST can address ambitious science goals ranging from the directdetection of hot exoplanets to the imaging of the heart of AGNs, addressing planetary,extragalactic and stellar science topics. All the science cases presented in this report canuniquely be addressed with CFHT giving the telescope an advantage over its competitors. Forexample, in the case of debris detection, it is shown that SPHERE cannot achieve the sameresult due to a limited dynamic range near the central object. The same conclusion applies forsome other topics. FIRST will therefore be very complementary to SPHERE.As far as time is concerned, we anticipate a 2.5 year building phase before first light whichcould happen, should the proposal be selected, in early 2012.


51. IntroductionWe present in this report the result of the feasibility study for a visible high dynamic rangeimager for CFHT.We first start with the presentation of the science cases:• Exoplanets• Debris disks• Young Stellar Objects• Cepheids• Novae• Compact nebulae around evolved stars• From stellar surfaces to mass loss• Active Galactic NucleiWe then present the principle of fibered aperture masking and describe the instrument conceptwe propose for FIRST. We show that the concept can be built given the current technology.We present the performance of the concept and discuss its compliance with the science goals.A laboratory demonstrator has been built during the feasibility study. First results arepresented.Developments necessary in the next detail design phase are presented.We give a preliminary description of the operations and data processing. The cost andmanpower needs are assessed in the last part of the report.


62 Scientific case2.1 ExoplanetsFIRST and extrasolar planets: contextFIRST open a new window in term of extremely high dynamic range at very low angulardistances. These are the typical characteristics which are required to make direct observationsof extrasolar planets. FIRST is expected to provide new limits in this domain, being hence oneof the instruments able to provide the first direct detections of the nearby extrasolar planetsorbiting close to their parent stars.Moreover, beyond challenging detections, FIRST wavelength coverage and spectroscopiccapabilities offer a unique possibility to scrutinize the atmospheric content of the observedextrasolar planet.Numerous models have been developed to predict the emergent spectrum from an extrasolarplanet from near UV to the infrared (Marley et al. 1999; Sudarsky al. 2003; Burrows etal. 2004; Hubeny et al. 2008; Fortney et al. 2008a). From these models a common pictureemerges in which the albedo and the spectral features in the reflected light from extrasolarplanets mainly depend upon the physical characteristics of the dust, condensates and haze inthe upper atmosphere of the extrasolar planets. Two classes of extrasolar planets have beenproposed: the pM class and the pL class. The class of planets that are warm enough to haveappreciable opacity due to TiO and VO gases have been named “pM class” planets, and thosethat are cooler have been named “pL class” planets (Fortney et al. 2008b). pM class planetshave temperature inversions (hot stratospheres) and present characteristics features of TiO andVO in the 600-800 nm wavelength range. The pL class planets absorb incident flux deeper inthe atmosphere where atmospheric dynamics will more readily redistribute absorbed energy.Most importantly for FIRST, we note that thermal emission from pM class planets in theoptical wavelength range is significant, making these planets attractive targets for opticaldetection (Fortney et al. 2008b).In case of detection with FIRST, even with a moderate spectral resolving power (dependingon the dynamic range achieved) the spectroscopic capabilities of FIRST will allowcharacterisation of the atmospheric content and structure. For instance, as already done in thecase of the transiting planet HD209458b, resolving the sodium NaI line at 589nm allows thecharacterisation of the atmospheric physical parameters (e.g., temperature, pressure) as afunction of the altitude (Sing et al. 2008a, 2008b). The potassium line at 770nm offers thesame opportunity in the visible wavelength range. At shorter wavelengths, variation of theRayleigh scattering as a function of wavelength depends upon the atmospheric variation of thepressure as a function of the altitude. Therefore the detection of the Rayleigh scattering allowsa direct determination of the temperature of the atmosphere at a few millibars pressurealtitude (Lecavelier 2008a, 2008b). Finally, we note also that TiO and VO molecules whichplay an important role in the thermal structure of the upper atmospheres are the mainabsorbers in the 620-800 nm wavelengths domain (Desert et al. 2008). Measurement of theemission spectrum of an extrasolar planet in this wavelength range, even at a resolving powerof R~10, will provide original information on the relation between thermal structure, albedoand condensates.


7As a conclusion, recent theoretical and observational results have shown that, even at lowresolution, spectra of extrasolar planets in the optical domain will provide useful constraintson the physical and chemical state of the upper atmosphere of these planets.Fig. 1. Planet-to-star flux ratios versus wavelength (in µm) from 0.5 to 3 µm for a one Jupitermass extrasolar planet orbiting a G2 V main-sequence star. This figure portrays ratio spectraas a function of orbital distance from 0.2 to 10 AU. For the 0.2 AU model, the temperatures ofthe atmosphere are high enough for the Na D doublet around 0.589 µm and the K I doubletnear 0.77 µm to be visible. These features are even more prominent for closer extrasolarplanets (the case in our study to reach lower contrasts). For cooler planets, the atmospherictemperatures are too low for the alkali metals to appear, but the methane features near 0.62,0.74, 0.81, and 0.89 µm come into their own. Water bands around 0.94, 1.15, 1.5, and 1.85µm (from Burrows et al., 2004, ApJ 609, 407).Observation of extrasolar planets with FIRSTUsing the published list of known extrasolar planets (~300 planets in September 2008), a firstlist of targets for FIRST can be established. Using the capabilities of FIRST as simulated in§5.10, we tabulated the planets for which a detection with a S/N larger than 2.5 can beobtained. We find that, from the list of presently known extrasolar planets, 6 extrasolarplanets can be detected with FIRST assuming a 10 hour exposure time per target. Since theplanets need to be bright they are all very close to their parent star. We have set a limit ofλ/4D for the closest distance possible. This is below the angular resolution but a fondamentalproperty of FIRST is that the PSF is accurately determined (by construction) to whitin thedynamic range. As a matter of fact, the PSF is set by construction by the array of filtered subpupilssince after filtering and after image reconstruction, the wavefront in the pupil is flat (towithin the noise) as well as the amplitude of incident lightwaves. A calibration star cantherefore be observed to model the PSF or to subtract the PSF from the science target (or


8better, the PSF can be reconstructed from the wavefront output by the reduction software).Since the dynamic range (or equivalently the noise) is constant in the image (the photon noiseis uniformly distributed across the image, this is the difference with classical imagingtechniques for which the photon noise depends upon the local intensity in the image), a faintcompanion with a relative flux brighter than the dynamic range will be detectable after PSFsubtraction. In principle we could have chosen to remove the constraint on distance but someconservatism is required here since the technique has not been tested on the sky. The λ/4Dlimit is derived from our experience with accurate interferometric instruments.This list can be extended when new extrasolar planets are detected.For the two best planets, because they require lower dynamic range, a spectrum could beobtained with a resolving power of R~10 (the exact value will depend upon the planetaryalbedo). These planets being low mass planets (GJ674b is a 12 Earth-mass planet, andHD69830b is 10 Earth-mass), the spectra information will be of prime importance. This willprovide crucial information: this will allow solving the question of their true nature and thediscrimination between Neptune-like or Super-Earth planets.Planet V mag Dist(pc)Star-Planet(AU)Star-Planet(mas)S/NdetectionS/NspectroGJ 674 b 9.4 4.5 0.04 7.1 9.2 2.9HD 69830 b 5.9 12.6 0.08 7.0 6.0 1.9HD 3651 b 5.8 11.0 0.28 9.9 4.2Gl 86 b 6.2 11.0 0.11 9.8 3.655 Cnc b 5.9 13.4 0.12 8.7 3.4HD 190360 c 5.7 15.9 0.13 8.2 3.0Table of FIRST exoplanet targets, and expected S/N ratio for detection and spectroscopy witha R=10 resolution power.2.2 Debris disksThe discovery of circumstellar disks is relatively recent since they were first observed aroundvery Young Stellar Objects (YSO) via radio observations of the so called “protoplanetary”disks which masses and sizes are in the solar mass and thousands of AU ranges. These disksare the environment where Jupiter like planets are formed during the first 10 million years. Itis in the 1980’s that a dust disk was really imaged for the first time around the star β Pictoris(Smith and Terrile, 1983) that revealed a new disk category, identified now as “debris disks”,much fainter, of the order of Earth masses and spread over less than 1000AU. These arewhere the telluric like planets do form in the next several hundreds of million years.These two types of disks are to be studied in detail to better understand the various andcomplex mechanisms which took place during the Solar system formation. FIRST will beparticularly adapted to the study of the “debris disks” which are extremely difficult to observebecause faint and close to their parent and shining star. Furthermore the regions of interest inthese disks are where planets form, i.e. in the few first AU away from the star, regions whichstay inaccessible to direct observations up to now. FIRST will have the capacity of observingmuch closer than before and particularly in the few AU region.


9In the following table we give the present list of 21 debris disks already imaged from eitherground based or space borne telescopes and visible from Mauna Kea (mainly Hubble):Name V RA DEC Distance(pc)Size(au)MassM-MoonE-EarthImaged Ref.#at (µm)ε Eridani 3.73 03 32 55.8442 -09 27 29.744 3.22 60 0.07 M 850 1α PsA 1.16 22 57 39.0465 -29 37 20.050 7.69 133-158 1.4 M 0.6 2Vega 0.03 18 36 56.3364 +38 47 01.291 7.76 80-120 0.5 M 70 3AU Mic 8.65 20 45 09.5318 -31 20 27.238 9.9 17-210 0.8 4HD 207129 5.58 21 48 15.7514 -47 18 13.014 15.6 0.6 5HD 10647 5.52 01 42 29.3157 -53 44 27.003 17.0 75-130 0.6 6HD 139664 4.64 15 41 11.3774 -44 39 40.338 17.5 60-109 0.6 7HD 109085 4.31 12 32 04.2270 -16 11 45.627 18.2 100-150 450 8β Pictoris 3.86 05 47 17.1 -51 03 59.5 19.3 30-1300 ~3 M 0.5 9HD 92945 8.67 10 43 28.2717 -29 03 51.421 21.6 55-170 0.6 10HD 107146 7.07 12 19 06.5015 +16 32 53.869 28.5 60-185 0.3 E 0.6 11HD 61005 8.22 07 35 47.4617 -32 12 14.044 34.6 240 1.1 12HD 15115 6.80 02 26 16.2447 +06 17 33.188 45.0 315E-550W 0.6 13HD 202917 8.67 21 20 49.9563 -53 02 03.144 45.9 0.6 5HD 181327 7.0 19 22 58.9429 -54 32 16.973 50.6 1.1 1449 Ceti 5.619 01 34 37.7788 -15 40 34.893 61.3 30-60+ 17.9 16HD 15745 7.50 02 32 55.8103 +37 20 01.045 64.0 -480 0.6 17HR 4796A 5.78 12 36 01.0316 -39 52 10.219 67.1 Ring 70 1.2 18HD 141569 7.0 15 49 57.7489 -03 55 16.360 99.0 500 0.5 19HD 32297 8.13 05 02 27.4366 +07 27 39.681 112.0 50-1680


10This table reveals the diversity of debris disks accessible to the FIRST observations whichwill be able to trace the dust distribution much closer than ever before, and in particulardown to about 1AU where precisely Earth like planets are presently forming within thesedisks.The above figure, pictures planetesimal belts around Beta Pictoris as indirectly inferredfrom the infrared observations (Okamoto, 2004, Nature 431, 660). These belts will bedirectly detected by FISRT. In this Beta Pictoris system, the inner belt at 6 AU is locatedat 300 mas from the central star and will be easily directly detected by FISRT. Any otherbelt closer to the star will be seen by FIRST at a distance down to 1 AU, and optical depthdown to ~10 -4 in one hour exposure time.But debris disks also contain gas produced like dust via collisioning and evaporatingbodies as originally suggested by Lecavelier Des Etangs et al. (1996). In the case of theparticularly well studied case of the β Pictoris disk (see e.g. a review by Vidal-Madjar etal., 1998), not only gas under many atomic (neutral and ionized) and molecular specieswere detected within the disk but even a still unknown braking gas able to stop the drift ofseveral of these species under radiation pressure has to be present. In particular we knowthis breaking gas cannot be neither HI nor H 2 (Lecavelier des Etangs et al., 2001).The gas in such disks could be a remarkable tracer of the dynamics going on as directlyshown by the localisation of neutral (e.g. NaI) as well as ionized species (e.g. Ca II) whichare seen to be in very different places within the disk (Olofsson, Brandeker, & Liseau,2001; Brandeker et al., 2004) although shown to rotate in keplerian motions as far asabout 100AU away from the star.This studied unique case of β Pictoris has to be extended by FIRST observations at muchcloser distances from the star where we know the action takes place, as well as over otherdebris disks where we already know that gas has been detected. These are listed in thefollowing table:51 Ophiuchi Roberge et al. (2002), Dent et al. (2005)σ Herculis Chen & Jura (2003)HD 32297 Redfield (2007)HD 158352 Roberge & Weinberger (2008)


11HD 118232 Roberge & Weinberger (2008)HD 21620 Roberge & Weinberger (2008)HD 142926 Roberge & Weinberger (2008)HR 333 Chen et al. (2006)49 Ceti Chen et al. (2006), Dent et al. (2005)This list could be extended by the following table where the disk in which the COmolecule was detected, taken from Dent et al. (2005):The outside radii of all these CO detections are in the range of 0.4 to few arcsec showingthat the detailed observations of the structures of all these disks are clearly within thereach of FIRST provided that their gas content in other gas species will be high enough, avery probable situation since these are all similar or denser than β Pictoris.Again a relatively large sample of debris disks will be accessible to FIRST observationswhich should reveal possibly how Earth like planets were or are presently formed withinthese “second generation” disks.One question could be raised about the possible performances of SPHERE (Beuzit et al.,2006) over the same scientific objectives. Taking the specifications mentioned inBoccaletti et al. (2008), related to a 1.68 µm wavelength range: “In the H band, thecontrast reaches 10 -4 at 2.5 λ/D and about 10 -5 at 6λ/D. Performances are slightlydegraded at shorter wavelengths since the impact of phase aberrations are more critical.”Boccaletti et al. (2008) conclude that the SPHERE coronographic performances are: “Theoff-axis transmission (combining mask attenuation and stop throughput) reaches 0.58 at0.1”and 0.73 at 0.5”. ”These performances show that the two instruments, SPHERE and FIRST will operate intotally different distance domains from the central star, the FIRST investigations givingaccess to the first AU away from the star, i.e. where we expect to see directly theplanetary forrmation process in action, while the SPHERE observations could only give


12access to much larger distance ranges where only the indirect consequences of theseprocesses could be seen.FIRST is thus the best possible instrument for direct observation of planetary formationprocesses.2.3 Young Stellar ObjectsThere is now strong evidence that in many low to intermediate young stars the circumstellarmaterial is in a disk-like structure from direct imaging with millimeter interferometry (e.g.Mannings et al. 1997, Nature, 388, 555), infrared imaging (e.g. Grady et al. 2005 ApJ, 630,958), and evidence for rotating disks based on the shape of spectral lines (e.g. Acke et al.2005, A&A 436, 209). A large fraction of these young stars are associated with jets and/orbipolar outflows. They seem to be an essential ingredient of the star-forming process becausethey are believed to be the way of carrying away the excess angular momentum from accretedmaterial onto a star. Despite their important role in star-formation, jet phenomena still remainsmystery.The origin of Jet launching and collimationThe most important question is how are jets launched in accretion disk system. It is not knownwhether jets originate from the interface between the star's magnetosphere and disk (X-windmodel, Shu et al. 1994, ApJ, 429, 781), or from a wide range of disk radii (disk wind model,Blandford & Payne 1982, MNRAS, 199, 883), or stellar winds channeled along field linesemerging radially from the star (stellar wind model, Romanova et al. ApJ, 635, L165). Thesemodels predict that jet excitation condition often have a size not larger than a few tens ofAUs. Therefore to discriminate among these models, one needs information on the smallestpossible spatial scales, at least 0.1” (=14 AU) even for the nearest star forming region, whichallows to probe the structure, excitation and kinematics of the jet close to its source.Fig. 1 The jet emanate from the source in a direction perpendicular to the dark lane representing thecircumstellar diskExtensive high-angular resolution studies of jets and outflows have been carried out withHST and ground based telescopes equipped with Adaptive Optics. A very interesting example


13is the HH30 jet in Taurus star-forming region (e.g. Ray et al. 1996 ApJ, 468, L103; Watson etal. AJ, 133, 845) Fig. 1 shows HH30 edge-on disk of dust around a young low-mass star.Light from the star illuminates the top and bottom surfaces of the disk, while the star is hiddenbehind the dust. The jet emanates from the inner region of the disk and it is well collimated,confined to a narrow beam. Another example is a micro jet from Classical T Tauri stars,obtained with the AO system on CFHT (Fig. 2, Dougados et al. 2002, Rev. Mex. Astron.Astrofis., 13, 43). One can see that a complex morphology for the jet emission close to thecentral source. From these observations are we have a good understanding of how theypropagate and interact with their surroundings on large scales. However even at the highestspatial resolution imaging currently available, the jet inner most region (below ~ 14 AU) areunresolved, where the jet collimation and acceleration occur. A fundamental difficulty inimaging a faint jet close to a bright star is high contrast with the source itself and a lack ofangular resolution.The angular resolution of a 3-meter class telescope equipped with FIRST will be higherthan 50 mas and the unprecedented contrast better than 10 3 can be achieved. It willallow us to observe the jet launching region directly and it will provide the moststringent constraints for the various jet launching models.Fig. 2 [OI]λ6300 + continuum narrow-band image of the DG Tau jet obtained with the AO system PUEOat CFHT. Adapted from Dougados et al. (2002, Rev. Mex. Astron. Astrofis., 13, 43).Disk variability, jet rotation and precessionAnother interesting study that will become possible, thanks to the high-angular resolution andhigh dynamic range of FIRST, is the search for signatures of jet rotation around the symmetryaxis. This topic remains to be developed further.Filter requirementsJets and outflows show line emission produced by radiative cooling in post-shock zones. Thisline emission process is dominated by forbidden transitions, e.g. [OI]λ6300, 6363, [SII]λ6716,6731, and [N II] 6583. We therefore request the following filters:


14Filter λ 0 (nm) Δλ (nm)Continuum 600 60O I 630 2S II 671 2N II 658 22.4 CepheidsCepheids are the most famous variable stars thanks to the relation linking their pulsationperiod and their intrinsic luminosity (the P-L relation), discovered exactly a century ago(Leavitt 1908). This remarkable property has turned these supergiants into standard candles tomeasure extragalactic distances. As shown by Freedman et al. (2001) for the HST Key Project(KP) to measure the Hubble constant, Cepheids can yield the most accurate extragalacticdistances up to 30 Mpc. Since the physical mechanism driving their pulsation is wellunderstood theoretically (the "kappa"-mechanism, see e.g. Gastine & Dintrans 2008 forreferences), it is generally considered that Cepheids are "no-thrill" pulsators, although thecalibration of the zero point of the P-L relation to 1% level is still an open question,.Figure 1. Scattered light from the circumstellar envelope of RS Pup in the visible (left,from NTT/EMMI) and its thermal emission at 24 µm (right, from Spitzer/MIPS).Until recently, only one Cepheid was known to be surrounded by a circumstellar envelope(CSE), the long-period southern variable RS Pup (Figure 1). The situation changed last yearwhen Spitzer observations of the nearby Cepheids δ Cep and SZ Tau revealed in the thermalinfrared that they also are surrounded by large circumstellar envelope (CSEs). In addition,interferometric observations (Kervella et al. 2006; Mérand et al. 2006; Mérand et al. 2007)revealed compact CSEs around four nearby Cepheids (L Car, Polaris, δ Cep, Y Oph), out offour observed stars at a scale of a few stellar radii, giving a 100% prevalence in the sample.These observations converge towards evidence that most Cepheids, if not all, are surroundedby dusty CSEs contributing significantly to their near and thermal infrared flux.


15The presence of CSEs complicates the Cepheid picture on two fronts: they impact theirapparent brightness, potentially introducing a bias on their P-L distances, and the evolution ofthe stars that reach the Cepheid spot in the HR diagram must now include very significantmass-loss. In other words, both the evolutionary status and standard candle status of Cepheidsare now in question.Light scattering properties of the Cepheid CSEsThe FIRST high contrast imaging of nearby Cepheids in the continuum will map the scatteredlight properties of the CSEs very close to the star. We will retrieve in particular the colordependentsurface brightness of the CSEs. Using a simple light scattering model, these datawill allow us to derive the physical properties of the envelopes (dust size distribution, density,total mass, composition, circumstellar reddening). The derived total mass in particular(assuming a standard gas to dust ratio) is the required input parameter for the evolutionarymodeling of Cepheids.There is a long-standing discrepancy between the evolutionary (derived from stellar evolutionmodels) and pulsational (from dynamical pulsation models) masses of Cepheids. The questionis still not settled today (see e.g. Caputo et al. 2005), and will greatly benefit from the"weighing" of Cepheid CSEs using FIRST (and supplemental) data.In addition to this evolutionary problem, our recent interferometric observations have shownthat the Cepheid CSEs may create a 5% bias on the near-infrared magnitudes of these stars(more in the thermal-IR domain), and thus on their Baade-Wesselink distances and position inthe P-L diagram. A refined knowledge of the properties of the CSEs in the visible is requiredto evaluate the impact of their presence on the apparent brightness of the stars.Cepheid distances from their light echoesFrom the monitoring of light echoes in the circumstellar nebula of RS Pup, we derived itsgeometric, reddening-free distance with a 1.4% accuracy (Kervella et al. 2008). At the time ofthese observations, this star was the only Cepheid known to be surrounded by a sufficientlylarge light-scattering nebula to allow such a measurement. We started in 2008 a program withthe CFHT/MegaCam CCD camera to obtain similar observations of the CSE of δ Cep andretrieve its distance geometrically to a few percent accuracy.But these observations concern the parts of the nebulae located far from the star, thanks to thefavorable ratio CSE size/distance of these two Cepheids. As shown by the interferometricobservations, Cepheid CSEs more commonly appear as compact envelopes. The extension ofthe promising "light echoes" method to a more significant number of stars will require theobservation of the faint scattered light nebulae very close to the Cepheids, within a few 100mas from the star. FIRST has the potential to image the compact CSEs in scattered lightthanks to its high-contrast imaging capabilities. This will enable us to apply the light echomethod to several additional Cepheids (possibly ~10) to retrieve their geometric distanceswith high accuracy.While we do not expect these few distances to be sufficient to calibrate the P-L relation to1%, they will serve as fiducials to derive the p-factor of Cepheids over their full range ofperiods. Fundamental for the Baade-Wesselink method, this period-dependent factor is


16currently the source of a ~4% systematic uncertainty, that will be relieved by the FIRST lightecho parallaxes. As the BW method is applicable to hundreds of Cepheids, this will provide arock solid 1% calibration of the P-L relation.Binary CepheidsAs demonstrated by several authors (Szabados), binarity and multiplicity are very common inthe Cepheid class. The Cepheid binaries with known periods are mostly 1-5 years (about 20 ofthem). FIRST has the required sensitivity to resolve the close-in Cepheid companions such asthe companion of Polaris (Evans et al. 2008, Figure 2), for which the contrast in the V band isaround 7.2 mag at a typical separation of 170 mas. The high contrast between the Cepheidsand their main sequence companions makes it a difficult observation, hence the need forFIRST. In order to progress on the question of the pulsational and evolutionary massdiscrepancy, precise astrometric masses are highly desirable.Figure 2. HST ACS images of the Polaris A and its companion in the UV (220 nm). Thecompanion is visible left and below Polaris A (figure from Evans et al. 2008).2.5 NovaeFIRST will be able to image in the visible the shape of bright explosions like Novae. Thehuge dynamic range and spatial resolution provided by this instrument is perfectly suited todetect the first hint of the extended ejecta around a very bright outbursting source (dynamicalrange required better than 10 5 at the very first moments, decreasing after).A classical nova eruption results from a thermonuclear runaway (TNR) on the surface of awhite dwarf (WD) that is accreting material from a companion star in a close binary system.In the explosion, a mass ranging between 10 -8 and 10 -5 solar mass is ejected at velocitiesbetween a few hundred and a few thousand kilometers per second.Most galactic novae remain spatially unresolved for several years, while the evidence of theirasymmetrical shape is increasing (as revealed for instance by polarimetric measurements).The highly time-pressured 8m telescopes are rarely used in this field. The recent impressiveimage of the outburst of the recurrent nova RS Oph is a good illustrative case. The ejectionwas highly asymmetrical but the radio images were limited by the physical appearance of thenebula in that wavelength regime: a bright ring. HST was used 150 days after the outburst, intwo narrow band filters centered on forbidden lines: the pictures revealed an impressivebipolar nebula (Bode et al. 2007). The central source was compact, and the nebula relativelyfaint: the HST dynamic range and spatial resolution hardly could monitor the source hereafter.


17These observations were crucial to ‘clean-out’ the picture of the event. Let us suppose anasymmetrical outburst of this kind involving polar/equatorial ejection velocities of3000/1000km.s -1 . The 3.6m CFHT has a spatial resolution in Hα (658.1nm) of ~47mas, andin Hβ (0.486 µm) of ~33mas. At 2 kpc, a bipolar structure can be studied typically 50-100days after outburst, given the high-dynamical range of FIRST. A rate of 4-8 novae per year,ranging from m v =3 to m v =14 are routinely detected and monitored by non-professionalastronomers worldwide. Optical interferometers can tackle the very first moments of thebrightest events, but the number of novae that will be studied with this technique should berelatively limited: 1-2 objects per year. For these events, a coordination with CHARA orVLTI observations should be possible. FIRST would represent a great complement,providing an accurate view of the nebula, soon after the sparse uv coverage observations fromoptical interferometers. In particular, the VEGA/CHARA combiner could study the very firstmoments in the Hα region, then would be followed by a subsequent FIRST monitoring withbetter imaging capability.FIRST alone could play a major role in this field by providing images for a statisticallyinteresting sample of events, allowing to determine the ‘shape-corrected’ distance parallaxes,but also to provide clues on the physical mechanism(s) at the origin of the asymmetry usingdifferent images from different lines at various excitation states. As shown in the figurebelow, the Field-Of-View to resolve these objects is not larger than 0.5”. Useful narrow bandfilters: Hα, [OIII]λ500.7, [NII]λ654.8.HST image of the ejecta around the recurrent nova RS Oph obtained about 150 days after theFeb.2006 outburst. A high dynamical range is necessary to separate the emission from theejecta and the outbursting source, and narrow band filters are necessary to isolate interestinglines. A bipolar shape is clearly visible in the deconvolved image (right).The list of filters required to achieve these observations is given below:Filter λ 0 (nm) Δλ nmHβ 486 2 to 3[OIII] 502 2 to 3[NII] 654.8 2 to 5Hα 658 2 to 6


182.6 Compact nebulae around evolved starsThe Hubble Space Telescope has been extensively used by the evolved star community. Inparticular, HST has provided numerous images of Planetary Nebulae with exquisite details.These images have deeply rejuvenated deeply the field as they revealed many phenomena atthe origin of the complexity of the nebula that were not considered before.Blue SupergiantsOne problem frequently encountered by HST was the brightness of the central source and thecontrast over the close stellar environment. It may be not such a surprise that the famousunstable supergiant η Car was so well studied by this instrument, since the dense bipolarnebula absorbs efficiently the central source, reducing naturally the contrast. η Car is a specialcase, and many other massive stars could not be studied so well. It is also worth mentioningthat the coronographic device has rarely been used.Luminous Blue Variables (LBVs), such as η Car are massive stars caught in a stage ofinstability, leading to huge eruptions that can be misinterpreted as supernovae. This stage isvery short, typically 10 4 yrs, and the number of LBVs is very small, a handful. However, manyblue supergiants exhibit compact nebulae in their surrounding that most likely witness pastviolent ejections.These objects, the so-called the B[e] stars, can easily be confused spectroscopically withLBVs, as they share many characteristics. But their main distinctive feature is that theyharbour a permanent environment of hot dust, coupled with a photometric stability in contrastwith the LBVs.The evolutionary link between the two stages is not defined yet, but as massive stars aremaking one big loop toward the red, and then blueward, in the course of their evolution,continuously losing mass at high rate, the blue supergiant region is filled by many spectraltypes having very different initial and current masses.Marston & McCollum recently searched for large scale nebulae around 25 B[e] stars, findingfor about half of them evidence of complex, ring-like or bipolar nebulae. For all these starsthe picture of the environment in the close vicinity, the source of the wind, is missing.For instance, a campaign has been undertaken on the mysterious system HD87643: NACOobservations showed that this source is resolved, but the limited Strehl ratio (~0.2-0.4) and thelimited dynamic range (a few 10 3 ) did not allow to recover the geometry of the source in theinner 1”. Such a high resolution study of these stars was never really performed in theNorthern Hemisphere. The study can be performed in broad band (FWHM~20-50nm) near0.6 or 0.7 micron, or in a narrow band (FWHM~2-6nm) filter centred on Hα. The dynamicalrange required to see the nebula depends on the distance from the bright star: typically 10 5 at0.1-0.2”.


19Fig.? Left: Recent B, V, R image of the B[e] HD87643 circumstellar environment at large scale (ESO/WFIinstrument), for which a dusty disk was discovered by MIDI (direction indicated by the red line). The detail inthe close vicinity at the arcsec scale is missing due to the saturation of the detector and the low resolution of theinstrument. Right: MWC137 and its complex environments (the arcs are at 70” from the star), in Hα. The coreis saturated (from Marston& McCummum, 2008).Compact environment around AGBs, post-AGBs and young Planetary NebulaeStars having a main sequence mass between 1 and 8 solar masses should become giants twicein their lifetime. The second giant phase is called the Asymptotic Giant Branch (AGB).During the AGB, these intermediate mass stars undergo severe mass-loss and the post-AGBphase is the rapid evolution of the star following the AGB giant sudden collapse. During thepost-AGB phase the stellar core rapidly contracts and heats, while the outer cool and dustylayers become disassociated from it. Eventually the central star (CS) becomes a white dwarf(WD), while the gas lost during the AGB and post-AGB phases becomes ionized and shinesas a planetary nebula (PN). The drastic changes observed in AGB/post-AGB starscircumstellar structure are particularly puzzling. During the late AGB or early post-AGBevolutionary stages, the geometry of the circumstellar material of the majority of stars,changes from roughly spherical to axially symmetric, or point-symmetric (multipolar). Thisvery diluted environment, almost decoupled from the star, acts like an amplification of thephenomena operating in the deeper layers, imprinting their mark in the shaping of the nebula.Most of the models proposed to explain the formation of the axisymmetric nebulae rely on thepresence of an equatorial disk that favours a polar mass-loss. If the disk is precessing orwarped then multipolar nebulae with ejection lobes in different directions can also beobserved.The pre-planetary nebulae are the most interesting targets. They are just beginning to formand ionize a dense, compact dusty nebulae, that seems carved by jets that shape the ionisedmaterial. These jets rapidly change their main direction, and transient warped disks have beeninvoked to explain this phenomenon. The VLT and the Keck observatories obtainedimpressive images in the near-IR, but without resolving the inner cores. The greatest potentialfor detecting these disk is with infrared interferometry, and in the visible, thanks to thescattered light. Continuum and Hα images with a field of 0.5’’ are necessary to conduct thisstudy. We note however that many of these targets are faint in the visible (in the range ofm V ~8-12). The Keck and the VLT in K and L are complementary in terms of spatialresolution compared with the CFHT in visible (see an example with the HST in Lagadec et al.2006), as they probe the hot dust, and FIRST essentially the close scattered light. Dedicated


20studies can be undertaken on about 20-40 targets, as a complement to the surveys currentlycarried out in the mid-IR (Gemini and VLT).A (Southern sky..) example: CPD-56°8032 (Chesneau et al. 2006). The HST image is saturated in the core,preventing study of the inner 0.5 arcsec region. STIS spectra are split into two parts, suggesting a disk seen athigh inclination.2.7 From stellar surfaces to mass lossFor most stars, the final phase of hydrogen burning occurs on the Asymptotic Giant Branch.The AGB stars expand until the atmospheres become very tenuous, the surface gravity low,and the star’s instability to pulsations leads to rapid mass loss in the Mira phase, returningnuclear products to the interstellar medium, and producing the extended remnants later visibleas planetary nebulae.Figure A shows a schematic diagram of a Mira-type star, according to current understandingof its structure. In this diagram, R i is the deepest layer that can be directly studied atwavelengths of low opacity (e.g. near 1.6 µm) and R o is the highest layer which is opticallythick at some wavelengths (e.g. in strong molecular bands). Direct measurements by opticalinterferometry (Perrin et al., 2004) show that typically, R o ~ 2R i . Depending on thewavelength of observation and the pulsational phase, the apparent radius of the star may beanywhere in approximately this same range, R i to R o . Correspondingly, by choice ofwavelength, the atmospheric content and dynamics can be probed as a function of depth.Somewhat further from the star, a ring of SiO masers is found (Cotton et al., 2006). Anexpanding envelope of dust and gas surrounds the star, detectable by emission from the dustcontent (Tatebe et al., 2006) and by Doppler shifted absorption from the molecularconstituents.


21Figure A. An idealized diagram of a Mira star and near circumstellar region.The real stars are certainly far more complex than this conceptual diagram. In the lowdensities of the Mira upper atmosphere, there is no certainty that the star is either spherical orhomogeneous, and substantial reason to expect that it is not (witness the exotic shapes ofplanetary nebulae).Imaging of these stars is still in an early stage of development. Both aperture masking(Woodruff et al. 2008) and interferometric array imaging have been employed. Figures B andC show two of the few examples of the latter.Figure B. Reconstructed interferometric images of the Mira R Aqr (Ragland et al. 2008) at1.51, 1.64, 1.78 µm.


22Figure C. Our reconstructed interferometric image of the Mira omicron Ceti (unpublished) inthe K band.Red supergiant stars have some structural similarities to Miras, with the deep, tenuousatmosphere, irregular, though with only low amplitude oscillations, but exhibiting maseractivity, extensive mass loss and dust shells. In these stars, convective motions dominatephotospheric energy transport, and 3-D hydrodynamic models give an indication of whatmight be observed with imagery. Figure D shows an example of predictions (A. Chiavassaand B. Freytag, private communication) for the appearance of a red supergiant in the K-band.Figure D. Appearance of the surface of a cool supergiant, integrated K-band, based onnumerical hydrodynamical calculations (A. Chiavassa and B. Freytag, private communication,2008).


23Figure E shows an observed image which illustrates the state of the art in confrontingobservation to model supergiants, with imaging from an interferometric array.Figure F shows the predicted appearance from the same model of the star in the visible-red.Owing to the stronger temperature dependence of the Planck function in the visible, thecontrast of the structure is much stronger.Figure E. Our image of the red supergiant, Betelgueuse, in the K-band, based onreconstruction from interferometric array measurements (unpublished).Figure F. Appearance of the surface of a cool supergiant, integrated 610-680 nm, based onnumerical hydrodynamical calculations (A. Chiavassa and B. Freytag, private communication,2008).


24While the correlation of the observed images to the conceptual model of Fig A is not exactlyknown, the observed images correspond to a region in the range of size about R o . The imagesappear to suggest some surface structure on the face of the stars, and significant asymmetry inthe outer layers of the atmosphere, without yet revealing any obvious pattern.In order to directly explore further the evolution and fate of the AGB stars and redsupergiants, it is necessary to improve the imagery (for detection of both radial andasymmetric structure) with improved wavelength coverage (for determination of brightnesstemperatures of surface and circumstellar regions), to improve the dynamic range (for studyof the low brightness regions at the top of and above the photosphere), and the angularresolution (to distinguish radial structure). FIRST can contribute to all of these requirements,offering visible wavelength imagery with high dynamic range at the diffraction limit of thetelescope.For bright, nearby Mira stars, the apparent diameter R o is typically in the range 50-70milliarcsec (Woodruff et al. 2008), which may be compared to λ/D of 30 milliarcsec for a3.6m aperture operating at 500 nm. The extension of the atmosphere is larger as earlierspeckle measurements have shown (Labeyrie et al. 1977), with the dust being produced in theupper layers of the atmosphere and visible at still larger distances. Thus FIRST on a mediumaperture telescope will have a resolution capability well matched to the scale of the stellarradius and of the atmosphere for AGB star red supergiants with radii at least 10 mas, a fewtens of objects in the northern hemispheres. It will be possible to produce continuous imagesof these star from the photosphere up to the inner part of the dust shell (a few 100 mas) thusgiving access to a full picture to study the mass loss process (in the visible the dust will betraced through scattered light).Complemented with additional longer interferometric baselines observed with existing arrays(VLTI, CHARA) it will be possible to reconstruct images that can directly test the predictionsof Figure F for supergiants, and to investigate the surfaces of Mira stars, which at present arecompletely unknown.Stellar surfaces and close environments can be explored with moderate spectral resolution. Abasic requirement is to observe in narrow band filters isolating Hα to detect shocks and tracepulsations in the atmosphere of the objects. The atmosphere of Mira stars contain a largequantity of molecules like TiO from the base of the atmosphere up to a few stellar radii. Thesecan be studied in filters defined by Scholz & Takeda (1987). In addition, the stellar surfacecan be measured in the continuum at 400 nm if the instrument allows such a short wavelength.A basic set of filters is defined in the following table:Filter λ 0 (nm) Δλ nmContinuum 400 5Hα 656 60TiO 712 20Pseudo continuum 754 20


252.8 Active galactic nucleiUnderstanding the structure of the close environment of Active Galactic Nuclei and the waythe central engine can be fueled with matter are two major issues in the study of nuclearactivity in the core of galaxies. Several pieces compose the puzzle: a ionization cone, anarrow line region (NLR), a compact starburst, a broad line region (BLR), a dusty moleculartorus and micro-spirals and micro-bars. Difficulty appears when imaging directly thenucleus: the high contrast between the core and the low surface brightness of its surroundingsis a critical problem. Then two effects avoid deep probing of the surrounding: the shortexposure time, to avoid saturation on the nucleus, do not allow to reach photon noise in itssurrounding; and the scattered light from the nucleus that pollutes significantly the extendedemission. These kinds of problems are common in other programs, for instance the search forfaint companions or disks around stars, where a high dynamical range is required.Coronographic observations of Active Galactic Nuclei (e.g. NGC 1068, Gratadour et al. 2007)reveal complex structures due to accretion and ejection process. In the near infrared, suchstructures usually trace thermal emission from dust, and some hydrogen (molecular orionized). In the visible, the view is much more complex, due to the very variable absorptionand higher optical depth, but also due to the complex chemical composition and morphology(e.g. OIII ejection cone). This is especially true for Seyfert 2 galaxies where the moleculartorus or clouds of the Unified model hide the central source and produce a complex andpatchy environment.However, the highest contrast objects, such as Seyfert 1 galaxies are expected to be much lessobscured and may in fact afford a view of the nucleus and its close environment, if thecomplex structural PSF artifacts can be removed or calibrated. The contrast ratios aresomewhat less challenging than in the case of exoplanets. Among the structures that can beobserved at an unprecedented angular resolution, given the gain in contrast are the starformation rings and molecular torus, the Narrow Line Region, the jet environment andindirectly the Broad Line Region. Of course the later is of dimension much too small(typically a few hundred AU) with respect to the achievable resolution (say 5 kpc) : howeverrecent studies (Gratadour, Rouan A&A accepted) showed that a large fraction of the lightemitted by the BLR could be reflected through Thomson scattering on the electrons of theNLR : this indeed is what is observed on NGC 7469.To capture this emission, however, an exquisite angular resolution is still needed becauseotherwise there would be a rapid dilution of the BLR emission by the strong flux of the starforming region. Many spectral lines could be used as different probes of the differentstructure: coronal lines, OIII, NII, Hydrogen recombination lines. As regards jets studies, afirst target would be the exploration in M87 of the structure of the jet at its very beginning, asclose as possible from the central engine. Another goal is BL Lac which are the extreme caseof a Seyfert 1 with the jet emitted by the nucleus being practically on the line of sight. This isanother class of object where the very close environment of the AGN could probably beexplored with FIRST, a situation that was not possible until now because of the strongcontrast. BL Lac which are known to exhibit a very flat spectrum because it’s dominated bythe synchrotron emission of the jet, could reveal the actual nature of their surroundings whoseemission could be spatially separated from the jet’s one.19 Seyfert 1 and 4 BL Lac with visible magnitude below 13 are observable in the northernhemisphere and would constitute the basic sample of this program (list below).


26Source name RA DEC V z Seyfert typeNGC 4594 12 39 59.3 -11 37 23 9.25 0.002 S1.9NGC 4395 12 25 48.9 +33 32 48 10.27 0.001 S1.8NGC 2985 09 50 22.1 +72 16 45 10.61 0.004 S1.9NGC 1068 02 42 40.7 -00 00 47 10.83 0.003 S1hNGC 3982 11 56 28.1 +55 07 30 11.70 0.003 S1.9NGC 5033 13 13 27.5 +36 35 38 12.03 0.003 S1.8NGC 1161 03 01 14.2 +44 53 50 12.08 0.007 S1.9NGC 4168 12 12 17.3 +13 12 18 12.13 0.008 S1.94U 0241+61 02 44 57.6 +62 28 06 12.19 0.045 S1.2NGC 4138 12 09 29.9 +43 41 06 12.24 0.003 S1.9NGC 4565 12 36 20.6 +25 59 11 12.43 0.004 S1.9NGC 4639 12 42 52.5 +13 15 25 12.72 0.001 S1.0NGC 788 02 01 06.5 -06 48 56 12.76 0.013 S1hNGC 1386 03 36 46.4 -35 59 57 12.84 0.002 S1i3C 273.0 12 29 06.7 +02 03 08 12.85 0.158 S1.0NPM1G+78.0053 12 55 07.8 +78 37 15 12.90 0.043 S1UGC 3478 06 32 47.5 +63 40 23 12.90 0.012 S1nNGC 4051 12 03 09.6 +44 31 53 12.92 0.002 S1nNGC 1365 03 33 36.4 -36 08 24 12.95 0.006 S1.8


27SpatialresolutionField ofviewWavelengthcoverageSpectralresolution /bandwidth3 Science requirementsExoplanetsDebrisdisks≥ 7 mas ≤ 45 mas(1 au @221 pc)none 0,3’’-18.4’’(for entiredisk)inner part(≤ 0,3’’)unknown589-800nm visiblerangeMinimum:10Goal: few100Continuum:R=50.Atomiclines :R=5000 to50000YoungStellarObjects≤ 50 masCepheids Novae Nebulae StellarSurfacesbetter thanHST ie 47mas@Vbetter thanHST ie47 mas@Vbetter thanHST ie 47mas@V≤ 30 mas≤ 500 mas few 100 mas ≤ 500 mas 500 mas few 100mas600-658nmContinuum:60 nmOI, SII,NII: 2 nmas large aspossible(continuum)wide bandAGNbetter thanHST ie 47mas@Vfew 100mas486-658nm 600-700nm 400-754nm visiblerangenarrowbandfilters: 2-6nmHα: 2-6 nm20-50 nm@ 0.6 and0.7 µmmV 5.7-9.4 0-9 10-19 2-10 3-14 8-12 3-14 (largeamplitudevariables)Dynamicrange# sources 6 (more inthe future)5-60 nm wide band10 5 ≥ 10 3 ≥ 10 3 ≥10 3 10 5 10 5 @ 0.1- 10 2 ≥10 30.2’’≥40 ≥ 10 TBC 30 4-8/year 100 few 10 1913


284 Principle of fibered aperture maskingThe principle of fibered aperture masking has been described in Perrin et al. (2006) andLacour et al. (2007).4.1 The issue of high dynamic range imagingThe dynamic range is defined as the ratio of the maximum peak intensity in an image to therms noise level. Very high dynamic ranges can be obtained thanks to long integration times ingeneral. However, depending on the nature of the noise, the dynamic range will change. Ifnoise is pure photon or detector shot noise then the dynamic range will increase as the squareroot of the integration time. In addition to these additive noises, if speckle noise becomesimportant then it will quickly be the dominant limitation for high dynamic range imaging neara bright source. This is especially true if the imager point spread function (PSF) is not stableas different speckle patterns will be detected in the science object image and in the calibratorobject image with no possibility to directly disentangle a real faint source and a residualspeckle. In addition, speckles also produce an extra luminous background whose intrinsicphoton noise will limit the dynamic range, the fundamental noise limit for exoplanet detectionwith large telescopes (Cavarroc et al. 2006). Calibrator images can be used a posteriori withdifferential techniques to improve the detectability of fainter sources but it is clearly better toimprove the stability of the PSF first.Very demanding science cases such as exoplanet detection therefore require speckle noise tobe eliminated from the images for the two reasons mentioned above: stability of the PSF andlower level of photon noise.Speckle noise is all the more important as the wavelength is short. It is in particular verydifficult to suppress at visible wavelengths.An elegant way to reduce speckle noise is to control the wavefront surface with an adaptiveoptics system. High dynamic ranges (few 10 4 ) have been obtained at infrared wavelengths at afew diffraction limit distance from the central object (Roddier et al. 1996). The wavefront issampled at a given spatial frequency. Higher spatial frequencies are not seen by the wavefrontsensor and higher orders of turbulent phase are therefore not corrected hence some residualspeckles. Near the central object, the dynamic range is still smaller due to the residualinstabilities of the PSF and to the larger photon because of the larger local intensity. It is inpractice difficult to obtain a high dynamic range near the diffraction limit of the telescope.The only facility available in a near term in the visible will be the SPHERE ZIMPOLinstrument (Beuzit et al. 2006). ZIMPOL will allow the detection of a faint companion near abright star with a dynamic range of 10 -8 in a 4-hour exposure between 600 and 900 nm. Suchperformance will be achieved in a long exposure and at a quite large distance from the centralobject (1”). Also, this performance requires the companion to be at least 30% polarized (andthe central object must not be polarized) as the principle of the detection is a differentialpolarization measurement.Therefore, no universal high dynamic range imager exists at visible wavelengths.


294.2 Aperture maskingVery high dynamic range imaging requires to minimize the number of speckles in the imageand their intensity. Adaptive optics is an excellent solution. Because of shorter wavelengths, alarge number of actuators is required. It is nowadays absolutely feasible for a 4 m telescopebut with however some limitations with respect to requirements for high dynamic rangeimaging near the central PSF. On the order of 200 actuators are required to achieve diffractionlimited imaging at 0.5 µm on CFHT. But achieving high dynamic range performances of thelevel of that of SPHERE or GPI would require at least 2000 actuators in the visible and veryshort integration times making the adaptive optics system far more difficult and expensive.An alternative is data post-processing. One technique is none as lucky imaging. It builds onthe idea that some occurrences of atmospheric wavefronts are quite flat and can be selectedamong a set of short exposures. This may be a solution for diffraction limited images butcannot lead to high dynamic ranges because of a very limited exposure time and because ofthe difficulty to calibrate the residuals of the wavefront since there is no reason why theyshould be stable with this selection process.Another family of techniques is speckle imaging or speckle interferometry (Labeyrie, 1970).The fundamental idea is that atmospheric turbulences are frozen in exposures shorter than theatmospheric turbulence coherence time and that the turbulence-free spatial information can berecovered thanks to a post-processing of the data. The post-processing is performed in thespatial frequency plane, the plane conjugated to the sky plane by the Fourier transform. Itconsists in calibrating the energy at each spatial frequency thanks to a similar observation ofan unresolved star used as a reference. If turbulence sequences were the same for thecalibrator star and for the science target then disturbances due to turbulence would be ideallyeliminated by this procedure. Unfortunately turbulence is not stationary and differences occurbetween the two observations. Another shortcoming of the method is that different zones onthe pupil telescope will contribute to the same spatial frequency (this is called redundancy)and therefore turbulence phases will be mixed thus decreasing the energy recorded at a givenspatial frequency. Since decorrelation between phases increases with the distance betweenpatches in the pupil plane, two different sets of patches (or subpupils) corresponding to asame spatial frequency will have less and less correlated phases thus leading to an increasingattenuation of increasing spatial frequencies: the image gets less contrasted on short spatialscales than it is on large spatial scales. This effect is detrimental to high dynamic range.Another version of speckle imaging is called aperture masking (Haniff et al. 1987). Theprinciple is the same as for speckle imaging with a mask placed in a pupil plane to select nonredundantsub-apertures. With this set-up, spatial frequencies are formed with a single pair ofsub-apertures from the pupil plane. As a consequence high spatial frequencies are notattenuated with respect to lower spatial frequencies and the dynamic range in reconstructedimages is the same at all scales. This technique has been successfully used on severaltelescopes. As an example, the spectacular result obtained with Keck at near-infraredwavelengths on the pinwheel nebula (Tuthill et al. 1999). This technique still has somedrawbacks however as far as dynamic range is concerned. First of all, only a small fraction ofthe pupil is used: each sub-aperture has a size of the order of the Fried parameter r 0 . Becauseof the non-redundancy requirement, only a limited number of sub-apertures can be used and asmall fraction of the total pupil area is used to collect photons (2% in the case of the Keckexperiment) as illustrated by the figure of Tuthill et al. (1999) below:


30Second, in order to collect enough photons, the size of sub-pupils is such that the phase variesacross them inducing some decorrelation. The amount of energy fluctuates at the level ofgiven spatial frequency causing a speckle noise. Because of these two effects, the dynamicrange of the images collected at Keck is limited to 200.4.3 Aperture masking with single-mode fibersThe problems of aperture masking can be solved using single-mode fibers (Perrin et al. 2006).First of all, single-mode fibers are perfect spatial filters: any lightwave injected in a singlemodefiber comes out with a flat wavefront. As a consequence, wavefronts are fully coherentafter filtering with single-mode fibers. The price to pay is that the number of photonscollected by a fiber fluctuates as it depends on the wavefront in the input. Lacour et al. (2007)have shown that this can be measured during data processing (see § 4.4). From a practicalpoint of view, the wavefront can be considered ideally flat at a single fiber level and also atthe full pupil level after cophasing during the post-processing of the data. Consequently, animportant source of dynamic range loss, speckle noise, is eliminated.Second, fibers are flexible optics and it is possible to use the full pupil in the input andredistribute the sub-pupils in the output in a non-redundant way. The full photon collectingcapability of the telescope is thus used. The principle of the system is represented in thesketch below:This set-up actually turns the telescope into a multiple-telescope interferometer. Each coupleof pupils defines a baselines for which an amplitude and a phase of a visibility are measuredand calibrated. An image is then reconstructed from the full set of visibilities. This is analternative way to perform imaging. The classical one is to get directly an image in the imageplane (the focal plane of the telescope). The alternative is to measure the spatial frequenciesof the object in the Fourier plane (the cross-correlation of the pupil plane) and reconstruct the


31object from its spatial frequencies. The two are equivalent through a simple Fourier transform.The procedure to calibrate the data and reconstruct images is detailed in the next section.4.4 Photometric calibration and Image reconstructionThe fundamental idea behind the data reduction algorithm detailed in Lacour et al. (2007) isto recover a perfect Optical Transfer Function (the distribution of energy of spatialfrequencies in the spatial frequency plane) or equivalently a perfect Point Spread Function.This is done with a self-calibration algorithm which allows a simultaneous reconstruction ofthe PSF and of the image of the object.Recovering a perfect OTF means measuring the number of photons or the intensity injected ineach fiber (and normalizing it to a common value for all sub-apertures) and measuring theturbulent phase of all sub-apertures to subtract it from the full aperture phase screen to recoverthe object phase. The effect of the algorithm on the measured phases in the case of a pointsource is represented on the sketch below:!The redundancy of the pupil is a major feature for the algorithm to work. Redundancy meansspatial frequencies are measured several times in different ways. This is important as thetechnique of the closure phase is used to measure and remove turbulent phases. Triangles ofspatial frequencies are formed to apply this technique, phases of visibilities are summed overeach triangle, the result being independent of any turbulent phase. Closure phases aretherefore observables only characteristic of the object. The impact of using closure phases todeduce phases is that the reconstructed object is centered in the field whatever the initial realposition of the object in the field. The classical drawback of closure phases is that a singlephase information is obtained from three initial and ill-measured phases. For N sub-apertures,N( N "1) 2 pairs (or spatial frequencies) can be formed but only ( N "1)( N " 2) 3closurephases can be measured. Thanks to redundancy, if the number of independent pairs is smallerthan the number of independent closure phases then all the phase information (to within a tiptilt)can be recovered and unaffected by turbulence. In practice, this still holds if the numberof parameters to constrain (N photons or intensities and N phases) ! is smaller than the number


32( )( N " 2) 3 closure phases if less thanN( N "1) 2of measurements (not necessarily N "1pairs are obtained with the beam combiner, see the case in practice of the CFHT in § 5.4).Images are reconstructed for each recorded frame (each frame duration defines the exposuretime). The condition for ! this is that each sub-aperture sees at least one ! photon per singleexposure (in the photon noise limited regime). This condition sets the sensitivity of theinstrument. Reconstructed images can then be stacked (the total time is the integration time)to increase the SNR in the final image. An important characteristic of the reconstructedimages is that the noise statistics are uniform across the image unlike classical imaging.This is because the noisy measured quantities are the frequency components and that the noiseon frequency components is uniform. After Fourier transforming back to the image plane (theaction of the image reconstruction algorithm), noise remains spatially uniform.4.5 Theoretical performanceThe performance of the instrument is characterized by its sensitivity and by the dynamicrange of the reconstructed images as a function of the magnitude of the object.4.5.1 SensitivityAs stated before, the sensivity is set by the condition that at least one photon is detected perframe and per sub-aperture in the case of photon noise limited detection. The sensitivity thendepends on the size of sub-apertures. Lacour (PhD, 2007) has shown that the maximum sizefor a sub-aperture is 3r 0 for Kolmogorov turbulence as the amount of injected intensity doesnot increase beyond this sub-aperture diameter. As only a fraction of the light can be injectedin a single-mode fiber (mode matching), the effective size of a sub-aperture is indeed r 0 . For a2r 0 aperture, the effective size of the sub-aperture is still of the order of r 0 .4.5.2 Dynamic rangeBaldwin & Haniff (2002) have shown that the dynamic range of a final image reconstructedwith an interferometer is:DR =N pairs("V V ) 2 + "# 2with N pairs the number of spatial frequency measurements and δV and δϕ are respectively theerrors on visibility modulus and phase.!The typical accuracies achievable on visibilities are "V V = "# = 0.01 and the dynamic rangeafter N exp exposures scales as:DR " 35 # N! subN expthe equality only holds if all pairs are formed. In the case of 100 sub-pupils and assuming allpairs are measured then a dynamic range of 3.5x10 5 is achieved in 10000 exposures. Lacour etal. (2007) showed that this ! formula is valid in the case of FIRST. It means in particular thatthe dynamic range is photon noise limited and not speckle noise limited and that thereforespeckle noise has been eliminated. This particular formula needs to be scaled for a specificinstrument taking into account throughput, turbulence strength, source flux, etc …


334.5.3 Field of viewIn a single-mode interferometric instrument, the field of view is primarily determined by twothings:• the field covered by a single-mode fiber• the interferometric field of view i.e. the field over which fringes can be obtained.The first one is equal to λ/d with d the diameter of a sub-pupil. It therefore decreases with anincreasing sub-pupil diameter. It defines the maximum field of view accessible in a singlesnapshot.The second one is linked to the coherence length of the radiation. One can show that for aninterferometer with a baseline B working at a wavelength λ in a spectral band δλ, fringes canbe obtained in a field equal to:"# = $2B%$ = $ B & RFor a given field of view and for given B and λ, this sets the bandwidth and/or the spectralresolution R. In order to keep the field of view maximum, the spectral resolution is adjustedso that the field of view matches!the diffraction limit of each sub-pupil λ/d.


345 Instrument conceptThe concept of the instrument is very close to the theoretical view explained in the previoussection. Here we detail the main characteristics of FIRST at CFHT and discuss its subsystems.We show that the critical components have been identified during the study and thatthe instrument concept we propose is therefore valid. All critical components are alreadybeing tested in a lab experiment for which the CFHT fundings have been useful to buy thedetector. More detailed designs (mechanical, optical and electronics) can be worked out afterthis feasibility study step.5.1 OverviewIn this section, the general architecture of the instrument is presented. The subsystems areintroduced and their usefulness is justified.The basic need of the instrument is a convenient focus from which the pupil can be imaged onan array of focusing optics to inject the light into the single-mode fibers. The instrument canbe set at the Cassegrain focus of the CFHT. This has the advantage that FIRST can benefitfrom the AO to increase its sensitivity. From a functional point of view, the instrument can bedescribed in sub-systems as detailed in the figure below:Telescope (with or without adaptive optics)Image the telescope pupil onto the instrument entrance pupilCompensation of differential chromatic refractionControl of individual beams for efficient focusing in fibersArray of focusing opticsSet of fibers to filter beams and remap the entrance pupil into a non-redundant exit pupilBeam combination to interfere the beamsDisperse the light at the required spectral resolution in a photometrically efficient wayPhoton counting array or low noise detector arrayFIRST samples the pupil of the telescope. As a consequence, its entrance pupil needs to beconjugated to the telescope pupil thanks to the Pupil Imager.The working wavelength range is between 400 and 900 nm. In this range, differentialrefraction can be as high as 2 arcsec depending on zenithal distance. A classic AtmosphericDispersion Compensator is needed to cancel the chromatic elongation of the image. BecauseFIRST will reconstruct images centered by construction in each spectral channel (the phase


35recovery algorithm makes use of closure phases, see § 4.4), the requirement on the ADC is toensure that the elongation of the image in the focal plane of each sub-pupil is a fraction of thediffraction limit angle λ/d with d the diameter of a sub-pupil.Each sub-pupil has a size of a few r 0 . As a consequence, the wavefront is not flat and notperpendicular to the optical axis in each sub-pupil. Those fluctuations will not be corrected byFIRST but will be filtered out by the single-mode fibers. However, a steering system isrequired for each sub-pupil to bring each focused beam onto the corresponding fiber head tocompensate for static aberrations in the system (telescope and FIRST). Static aberrations willbe simple wavefront tilts at sub-pupil scale. An accuracy of a fraction of λ/d is necessary forthe Multiple Beam Steering System to avoid losing light on unresolved objects and to ensureaccurate pointings in fields larger than λ/d. The MBSS allows in addition to relax constrainson the accurate positioning of fibers at the focus of the instrument.Each sub-pupil needs to be focused into the corresponding single-mode fiber. This is the taskof the Multiple Beam Focusing System. Since as much as possible light must be collected bythis device, focusing elements must be placed next to each other. Besides, focusing in fibers isoptimum for a given beam aperture that depends on fiber characteristics (typically f-ratios of4-5 in the visible). In addition to these requirements, the optical quality needs to be high(Strehl ratio of 90%) for efficient beam coupling.One of the difficulties of FIRST is to simultaneously and efficiently inject light into all fibers.This requires fibers to be accurately aligned at the focus of the MBFS. They are a prioridifferent ways to achieve this. However, we have found, and this is explained in § 5.3, that theeasiest way is to combine an accurate MBSS with a Fiber Bundle where fibers are arrangedin an accurate fixed pattern matching the MBFS geometry (or of a scaled geometry). Thisgreatly relaxes the constraints on the fiber bundle (an accuracy of 1 fiber core is enough) anddoes not require each fiber to be adjustable thus making the instrument easy to align, tomaintain in the long term and to calibrate. The output of the fiber bundle can have anygeometry adapted to the beam combiner. This is one of the the advantages to use fibers: theoutput and the input can be decoupled and, in practice, the fiber bundle will be in two partslinked by connectors allowing to easily change the way beams are combined withoutchanging the injection part of the instrument. The fiber lengths need to be long enough toguide light to the beam combiner input and can be adjusted to make sure the whole devicedoes not introduce differential dispersion which would decrease the contrast of fringes.The conclusion we draw after this study is that the Beam Combiner remains the most difficultpart of the instrument. We propose a baseline solution which has been readily tested in the labduring the study (see § 5.12). It leads to a reduced number of combinations although theyprovide enough informations to reconstruct images with FIRST. A larger number of beamcombinations may however improve the image dynamic range and we outline some ideas toupgrade the instrument in § 5.9.For astrophysical reasons (spectral analysis) as well as for instrumental reasons (field ofview), a minimum spectral resolution is required. It can be achieved either with narrow bandfilters or with a Spectrograph. As a spectral resolution of 100 is required, this part of theinstrument can be based on a simple design with a prism. The more difficult part is the opticalsystem to perform the anamorphosis of the interferometric pattern in the direction of thefringes. A solution is being tested in the lab (§ 5.12).


36The last subsystem is the Detection System. We have chosen two classes of arrays: photoncounting detectors and EMCCDs which both allow to work in the photon noise limitedregime.The figure below displays a sketch of the instrument concept.A pupil imager is required to image the CFHT pupil on the 126-segment mirror, the input ofFIRST. A movable mirror can pick-up the beam from a source simulator for calibration andalignment of the instrument. A CCD camera has been inserted in the beam to control thealignment. Use of cube splitter is only presented here for the sake of simplicity but will bereplaced by an alternative solution in the final design of the instrument to optimize thethroughput of the science beam. The 126-segment mirror acts as individual steering mirrors toinject the beams in the fiber bundle through a lenslet array. The 126 fibers are recombined ina series of 14 1-D beam combiners each producing an anamorphosed image of the fringepattern dispersed with biprisms.5.2 Pupil sampling5.2.1 Sub-pupil sizesThe median seeing at CFHT is 0.65’’ at 0.55 µm wavelength corresponding to an r 0 of 15 cmin V. The maximum diameter of sub-pupils is d=3r 0 . This leads to a field of view ofλ/d=250 mas and an effective collecting area of r 0 . For nearly the same collecting area,choosing d=2r 0 leads to a field of view of 380 mas. We therefore set the baseline sub-pupildiameter to 30 cm.


37With a telescope diameter of 3.6 m and a central obscuration diameter of 1.58 m, the effectivenumber of sub-pupils is 116. In practice 126 sub-pupils will be required, some of them beingpartially field because of the hexagonal geometry of the lenslet array and of the segmentedmirror.5.2.2 Lenslet arrayTwo different options have been considered to sample the pupil. The first one is to build anarray of lenses. This was our initial choice as it would allow some space where to put thefibers and the 3-axis mount to accurately control the fiber position (the required accuracy is afraction of the fiber core – 4 µm – or, equivalently, of their sky angular size, λ/d). This apriori easy solution has actually several drawbacks:- a large number of actuators is required to control the location of the fibers;- the size of the system would be large: about 30 cm for the array of lenses thusmaking the global size of the instrument large;- no lens array is produced by the industry;- using individual off-the-shelf lenses to build the array would leave large gapsbetween the lenses and would be photometrically poorly efficient.Our preferred solution is to use micro-lens arrays produced for telecom applications. We haveused arrays from Süss Optics for our lab experiment (§ 5.12). The same company provideslarger arrays as the one we would need for FIRST. The array characteristics are as follows:µ-lens array from SüssMicroOptics• Array geometry: hexagonal• Lens geometry: circular• Material: fused silica• Filling factor (circle vs. hexagon): 70%• Lens pitch: 250 µm• Pitch accuracy: 0.25 µm• Focal length: 940 µm• Throughtput across the visible range including Fresnellosses (no anti-reflection coating): more than 90%• Image quality (Strehl ratio): 95 %• Numerical aperture: 0.12The Strehl ratio has been computed using Zemax. It leads to an injection efficiency of 78% inthe fibers.The pitch accuracy is high enough that it can be neglected in the following as it correspondsto 1/20 th of fiber core.As well as other parameters, numerical aperture is a critical one. As a matter of fact, singlemodefibers have their own numerical aperture and it has to be matched by the focusing opticsto avoid losing light. It is the case with a numerical aperture of 0.12 for the lenses as singlemodefibers are available with this same numerical aperture in the visible range.Taking into account the injection efficiency, the throughput and the filling factor of the lenses,the effective surface of each hexagonal sub-pupil is 50% the area of an initial subpupil. Thefinal effective surface of each sub-pupil taking into account micro-lens array throughput andimaging quality and the average corrugations due to turbulence at 0.55 µm is of the order of0.50 x π(r 0 /2) 2 . This is the number to be used in sensitivity computations (§ 5.10).


385.3 Injection in fibersInjection in fibers was the key issue at the start of this feasibility study. As a matter of fact,single-mode fibers are very sensitive to wavefront quality and to misalignment. As a rule ofthumb, the injected flux drops as a function of differential location between focused spot andfiber core according to a Gaussian law with the FHWM being equal to 0.5λ/d in sky angularscale units or half a fiber core. The requirement on fiber positioning accuracy is therefore veryhigh and typically of 1/5 th of the core or 1 µm.They are different ways to achieve this:• make a fiber bundle with a 1 µm accuracy pattern;• have a 3-axis actuator (xyz) for each fiber;• control the tip-tilt of individual sub-pupils to conjugate the fiber heads and the spots.At the beginning of the study, our preferred solution was the second one as the other two weretwo demanding. Although this is not difficult to do technology-wise speaking using piezoactuators for example developed at LAOG in France, this solution has an impact on thesystem (large number of subsystems) and on the space envelope of the instrument (a largespacing between the fibers implies a large focal distance for the sub-pupil array and thereforea large size for the individual sub-pupils). We therefore decided to look for a more compactsolution.Our preferred solution is to combine a precise array of fibers with a segmented micro-mirror.As detailed in § 5.5, the high accuracy of the fiber bundle pattern ensures an injectionefficiency of 37% for the according-to-specification performance and of 90% for themeasured performance for a flat input wavefront. Even without turbulence, the inputwavefront is not exactly flat because of static aberrations. The specification-compliantaccuracy on the fiber bundle does not ensure a sufficient injection efficient. For these tworeasons, the segmented mirror is necessary. It acts as a steering mirror for each sub-pupil andallows to reach injection efficiencies on a flat wavefront close to 100%. Since the stroke forthe mirror is not infinite, a mechanically accurate mounting of the segmented mirror, microlensarray and fiber bundle is necessary. It is achieved by mounting these on translation androtation stages.Thanks to this choice, the instrument is compact and the injection efficiency is primarilylimited by atmospheric turbulence.5.4 Beam combinationIdeally, all possible combinations should be measured by the instrument to have a maximumof redundancy to constrain phases and intensities in the pupil. This would be the game in aworld without photon noise. Increasing the number of combinations actually decreases thenumber of photons per recombined baseline. There is therefore an optimum and the number ofcombinations to achieve is not necessarily equal to:N pairs= N ( N "1)sub sub2In practice, with an instrument like FIRST for the CFHT, we are currently limited by asolution to combine as many beams as possible. As a matter of fact, combining more than 50beams non-redundantly is a challenge. A solution exists with a 2D combination scheme for 36beams. The drawback is however ! that the resulting image cannot be dispersed in an easy way


39thus preventing high spectral resolution on a continuous spectrum. We have therefore decidedto combine subsets of beams with a non-redundant1D arrangement. Accurate positioning offibers is obtained thanks to V-groove chips. The maximum of beams to be combined isactually limited by the maximum size of high precision V-grooves: 48. A 48-fiber V-grooveallows the non-redundant combination of 9 fibers. The locations of fibers in the V-groove aredisplayed on the figure below:This combination schemes produces fringes stretched in the direction perpendicular to the V-groove:Information is only carried across the fringe pattern. As a consequence, no information is lostif the fringes are collapsed perpendicularly to the V-groove. This operation is called beamanamorphosis and can be obtained by a combination of two cylindrical mirrors or anysomewhat close optical set-up. In this case, the whole interferogram can match a single line ofpixels of a detector. The interferogram can then be dispersed along the other axis thusallowing to obtain simultaneously high angular resolution and spectral informations on theastrophysical target.In the case of FIRST at CFHT, beams would be recombined in 14 subsets of nine fibers. Thetotal number of pairs is therefore:N pairs=14 " 9 " ( 9 #1)= 5042a number still much larger than the number of sub-pupils meaning that all the informationnecessary to reconstruct images free of turbulence will be available (there are 126x2unknowns to be determined).!We however would prefer a solution with a smaller number of subsets. As a matter of fact, theimpact on detection is not negligible with at least 2 pixels per fringe to satisfy the Nyquistcriterion. This particular point of beam combination is the one for which another solutioncould be studied to improve the baseline solution we propose which, we insist, guarantees thatthe instrument will work.Imaging of spectrally dispersed fringe patternsAfter anamorphosis, dispersed fringe pattern will look like this:


40Let us assume a 512x512 detector is used.Adequate sampling of each V-groove interferogram requires at least 2 pixels for the highestfrequency fringe. With a 48 V-groove the ratio in frequency between the lowest and thehighest frequency fringe is 44 (1 vs. 44). The closest spacing between afocal beams (producedwith V-groove fiber ends at the focus of a linear µlens array) is zero producing two fringesinside a fiber lobe imaged on the detector. The minimum number of pixels is thus 2 x 88 =176 pixels across the fringe pattern. This minimum sampling is for the shortest wavelengthfringes. The fringe pattern will scale with wavelength. Assuming a wavelength range of 400-800 nm, 352 pixels are necessary for a single-Vgroove. If restricted to 500-750 nm, thenumber of pixels is 256. Whatever the wavelength range, the sampling could be optimized forthe shortest wavelength and restricted to 256 pixels meaning that the wings of the fringepattern would be blocked for the longest wavelength (beyond 1.5 λ min , less than 50% energyloss in the worst case) to force its width to 256 pixels. In this case, 7 detectors would berequired (2 V-grooves per detector).512 pixels being available in the dispersion direction, and half pixels being required for eachpolarization, current detectors allow a spectral resolution in a wide band of 250 (250 elementsacross the 400-900 nm range, R=250 being the average resolution).5.5 FibersA sketch of the fiber assembly is presented below:It is in three parts:• The fiber bundle or fiber array;• 126 short length fibers for pathlength and dispersion compensation;• V-groove-bound exit fibers.2 connections are required for the whole assembly. The short length fibers provide flexibilityin the manufacturing of the fiber bundle and of the V-groove connected fibers. Besides, itallows to perfectly balance pathlengths and reduce dispersion which would otherwisedecrease the fringe contrast and the instrument sensitivity.


415.5.1 Visible wavelength range single-mode fibersFibers can be produced in different materials in the visible range: silica, plastic, fluoride glass.For astronomical applications, plastic must be avoided as no good single-mode fibers can beproduced with this type of material. The use of fluoride glass is possible but these fibers aremore difficult to handle than others as they are hygroscopic. Fluoride glasses having a verylow berifringence, polarization maintaining fibers are difficult to make with this material.Besides, they are very expensive. Silica fibers do not suffer from any of these drawbacks andare therefore our preferred solution. Besides, they are commercially available at reasonablecost.One of the key characteristics of single-mode fibers is their throughput. Losses may haveseveral origins:• Absorption by the material• Absorption by impurities• Losses due to micro-cracks• Losses due to scatteringThe latter effect (Rayleigh scattering) is the one that limits the length of fibers to be used atthe lower end of the visible band. In the visible range of wavelengths, this effect will inducelosses of 50 dB/km. In the case of first the fiber length will be of a few meters and thereforethe intrinsic loss in fibers can be considered negligible.We will use polarization maintaining fibers to ensure systematic maximum fringe contrast andwe will make use of FC/PC connectors which very low losses (0.3 dB per connection) and arewell suited to guarantee an alignment of polarization axes better than 3°. These specificationsare basic commercial specifications and do not require any custom components or dedicateddevelopment.5.5.2 Input fiber bundleAdequate fiber bundles are manufactured by Fiberguide Industries. Their characteristics are asfollows:• bundle pattern symmetry : hexagonal• fiber pitch: 250µm• pitch tolerance: 3µm (with a goal of 0.5µm) 1 µm accuracy has beenmeasured for the 36-element lab experiment sample (§ 5.12)• Flatness: 2µm• Connectors: Array to FC’s (key aligned to slow axis of polarization)• Fiber length: 50cmThe left panel of the figure below is a close-up view of the fiber bundle head with polarizationmaintaining panda fibers. The right panel is a side-view of the whole system.


425.5.3 Output fiber assemblyVery high accuracy 1D alignment systems have been developed for the telecom industry.They are known as V-groove chips. V-grooves provide high accuracy fiber positioning andpointing. V-groove chips manufactured by OZ Optics have the following characteristics:• maximum number of grooves: 48• fiber pitch: 250 µm• fiber positioning accuracy: 0.5 µm• fiber tilt accuracy : better than 0.3°A sketch of V-grooves manufactured by OZ Optics is presented below:V-grooves are etched in a silicon chip. Fibers are kept in position in the V-grooves with aPyrex lid laying on top of the silicon chip. Fibers are fixed to the grooves by UV curingmaking the whole assembly stiff. The V-grooves can accommodate any type of fibers as longas their diameter is standard (cladding diameter of 125 µm).The V-groove accuracy of the fiber spacing is of the same order as the accuracy of theconcentricity of the cores and claddings of fibers manufactured for the visible range (1 µm,FiberCore specification). The accuracy on core distances in the V-groove assembly istherefore of the order of 1.1 µm. This leads to an overlap between lightwaves from two fibersin the focal plane of the beam combiner of 87% and therefore of loss of contrast not higherthan 13%.It is to be noticed that the pointing angle of the fiber can be considered perfectly aligned withthe optical axis. As a matter of fact, the combination is performed in an image plane. Whatmatters is therefore the overlap between lightwaves in this plane. A pointing error of fiberswill produce a shift of lightwave locations in the beam combiner focusing optics plane. Whatmatters then is not to lose light. This is the case given the tiny pointing errors of fibers.5.5.4 DispersionFibers are material media of propagation with a refraction index dependent upon wavelength.Differential (interferometers are sensitive to differential quantities) chromatic dispersionarises when the second order term in the difference between the refraction index laws of twofibers cannot be neglected. The zero optical path difference is then wavelength dependentmeaning that a wide-band interferogram gets spread and its contrast is decreased. With


43homogeneous fibers, this is the case for uneven lengths of fibers in a fibered interferometer.The interferometer is all the more sensitive to differential chromatic dispersion as thephotometric bandwidth is large.The classic method to balance chromatic dispersion in a short-length fiber interferometer (thecase of FIRST) is to polish fibers to accurately match their lengths (rather than theirdispersion). This is achieved by controlling the polishing with an interferometer. Polishing isstopped when the fringe contrast is maximum. With this method, fringe contrasts of 95% areobtainedIn the case of FIRST, the effects of chromatic dispersion will still be more negligible. As amatter of fact, fringes will not be recorded in wide bands since a minimum of 12 in spectralresolution is needed to obtain maximum field of view.Last but not least, fiber lengths will be short enough and the instrument will be compact sothat the temperature of fibres can be considered uniform. No temperature induced dispersionvariation (due to a differential variation of fiber length with temperature) is to be feared.5.5.5 PolarizationFibers for the visible range are made of silica, a strongly birefringent material. Polarizationshave different propagation velocities and interferograms on the two polarization axes are notsynchronized causing a loss of contrast (100% loss if interferograms are shifted by half afringe). To overcome this issue it is necessary to use polarization maintaining fibers. Theseare commercially available. Interferograms obtained on the two axes of polarization need tobe detected separately. This is achieved using a Wollaston prism to detect the twopolarizations simultaneously. In the sensivity calculations, for non polarized sources, SNRneeds to be decreased by 2 . For polarized sources for which data on the two polarizationaxes cannot be combined, the SNR has to be decreased by a factor of 2.5.5.6 Throughput!With the short lengths of fibers used by FIRST, the intrinsic transmission of fibers can beconsidered 100%. Losses will be caused by connections (0.3 dB per FC/PC connector) andFresnel reflections (4% per reflexion). The fiber assembly will require 2 connections hence0.6 dB loss equivalent to a 87% throughput. Fresnel reflections will occur twice, at the inputand output of the fiber assembly. The total throughput is therefore 80%.5.6 Segmented mirrorTwo companies have developed MEMs segmented mirrors: Boston MicromachinesCorporation and IRIS AO Inc. This domain is quickly evolving and still better products areannounced for the coming year.Apart from functionalities, one big advantage of MEMs is that they are produced using semiconductorindustry techniques making the price quite low when compared to other types ofactive mirrors.We have investigated the products of the two companies for our lab experiment. BostonMicromachines was already able to provide a mirror with 1000 actuators. However, the strokeof their mirrors is quite limited: up to 5.5 µm but only 1.5 µm for their 32x32 actuator arraywhich is not enough to compensate for static aberrations. Their 12x12 actuator array howeverhas a stroke per actuator of 5.5 µm.


44At the time of our investigation, Iris AO could provide a 37 segment hexagonal mirror with abetter stroke: 5 or 12 µm. We are currently using a 5 µm mirror stroke which can be upgradedto 12 µm (low cost mirror exchange). Each segment has three actuators that allow to tip andtilt the mirror (differential voltage applied to the electrostatic actuators) and to apply a pistoncorrection (same voltage applied to the three actuators). This exactly meets our needs to steerthe beam in each sub-pupil of FIRST. The tilt range with a 5 µm stroke per actuator mirror is2.5 mrad at minimum (the tilt value depends on pointing direction and can reach 5 µm in thebest case). Given that each segment has a diameter of 700 µm and that the diameter ofindividual pupils for the lenslet array is 250 µm, the beam has to be reduced by a factor7/2.5~3 meaning a tilt angle produced by a mirror segment is magnified by a factor 3 at thefocus of a micro-lens. The minimum effective tilt range to be considered for a mirror segmentis therefore 7.5 mrad. With a micro-lens focus of 940 µm, the linear stroke at the focus ofmicro-lenses for a 5 µm stroke mirror is ±7 µm ie ±3 times a fiber radius. It scales to ±17 µmie ±7 times a fiber radius. On the sky this translates respectively to 1 arcsec and 2.3 arcsec.The 5 µm stroke mirror may be relevant for FIRST but it is probably better to choose the12 µm mirror to get enough margin. In both cases, our requirements are met.In addition, segmented mirrors from IRIS AO have excellent optical qualities: better than20 nm rms.The mirror we plan to use for FIRST is a 163 segment array. It is currently underdevelopment with a first sample available soon. The mirror geometry is presented below aswell as a photograph of the actuator array before the segments are deposited (in agreementwith IRIS AO Inc.). The red-colored segments are located in the central obscuration of theCFHT and will not be used. 126 segments are left.The diameter of the mirror is 8.2 mm.5.7 Differential refraction compensationCompensation of differential refraction in the visible band is a classical problem. We plan touse a system comparable to the ADC built for ESPADONS with two null deflection prisms.The performance of such a device is illustrated in the following figure (ESPADONSdocument # LAT 2 -ESP-TM-OPT-20):


45The residual of differential refraction after correction is 0.2 arcsec i.e. less than the field ofview of FIRST for a maximum zenithal distance.5.8 DetectionWe are considering two types of detectors:- Electron Multiplying CCDs (EMCCDs), a recent type of CCDs optimized for lowlight levels- photon counting detectors.5.8.1 EMCCD(description based on http://www.emccd.com/what_is_emccd/)EMCCDs are capable of detecting single photon events whilst maintaining high QuantumEfficiency, achievable by way of a unique electron multiplying structure built into the sensor.Unlike a conventional CCD, an EMCCD is not limited by the readout noise of the outputamplifier, even when operated at high readout speeds. This is achieved by adding a solid stateElectron Multiplying (EM) register to the end of the normal serial register; this register allowsweak signals to be multiplied before any readout noise is added by the output amplifier, hencerendering the read noise negligible. The EM register has several hundred stages that usehigher than normal clock voltages. As charge is transferred through each stage thephenomenon of Impact Ionization is utilized to produce secondary electrons, and hence EMgain. When this is done over several hundred stages, the resultant gain can be (software)controlled from unity to hundreds or even thousands of times.


46We have tested and used for our lab experiment an EMCCD bought with the funds providedby CFHT for this study. It is the HAMAMATSU C9100-13 whose characteristics are thefollowing:• Array size: 512x512• Pixel size: 16 µm• Frame rate: 31.9 frames/s• Exposure time: 10 ms (external synchronization)• RON: < 1e -• QE @ 600 nm: 92%The quantum efficiency across the visible range is displayed below:The QE is excellent in the 500-700 nm range and still reasonable until 800 nm but dropsbelow 400 nm and beyond 850 nm. Detection is photon noise limited across the entire visiblerange.5.8.2 Custom Photon counting detectorCPNG (Caméra à Comptage de Photons de Nouvelle Génération) is a camera built jointly byCentre de Recherche Astronomique de Lyon and Observatoire de la Côte d’Azur / FIZEAU tomeet the requirements of speckle interferometry: photon-counting fast read-out detection.It is described in : Rondeau, X.; Thiébaut, E.; Blazit, A.; Foy, R. “Performances of CPNG: anew generation photon-counting camera with real-time dedicated optimal processing”, 2006,SPIE, 6276, 62760L.CPNG is an Intensified CCD (ICCD) which uses electron multiplication in micro-channelplates to overcome the readout noise of fast CCDs. Amplification stages are made of astacking of evacuated intensifier tubes, each one made of a photocathode, a micro-channelplate and a phosphor screen. At the input of the first tube, the incoming photon is convertedinto an electron by a photocathode thanks to Einstein photoelectric effect. After thisphotoelectric conversion, an electric field draws the electron toward a micro-channel plate(MCP), where amplification occurs thanks to the secondary emission of electrons. Theacceleration of electrons is tuned via high voltages across the MCP. Bursts of acceleratedelectrons impact a phosphor screen at the output, which finally gives birth to thousands ofoutcoming photons from a single incoming one. Amplified photons are focused on a CCD forreading. The photocathode of the first intensifier sets the quantum efficiency of the ICCD.Two different types of photocathodes are required to cover the visible range:


47Characteristics of the detector:• Frame rate: 262 Hz• Maximum Quantum Efficiency: 36% @ 0.5 µm• Array size: 2000x2000 effective pixels (initial DALSA CA-D6 array: 532x516)• Pixel size: 9 µm5.8.3 Brief comparison of the two detector typesBoth detectors allow detection of photons in the photon noise-limited regime. TheHAMAMATSU detector has a higher Quantum Efficiency (factor of 2.5 to 7). CPNG has alarger number of pixels and a higher frame rate. The exposure time with the HAMAMTSUdetector can be reduced using a fast shutter but 10 ms seems enough for FIRST but should beconfirmed in the next phase of study. For sensitivity reasons, the baseline for FIRST is theHAMAMATSU C-9100-13 EMCCD.


485.9 Instrument upgrades5.9.1 Spectral resolutionIn this study, we only consider moderate spectral resolutions (limited by the detector size in awide band) of 250. For the science cases considered here, when necessary, higher resolutionscan be achieved through the use of narrow-band filters. It would however be possible toconsider a version of FIRST with a better spectral analysis capability at the cost of changingand upgrading the detection and spectroscopic systems. And at the cost of a poorer sensitivity.5.9.2 Field of viewThe field of view in a snapshot image (single integration) is primarily limited by the size ofthe subpupils and equal to λ/d. One possibility to increase the field of view is to decrease thesub-pupil size but at the cost of sensitivity. A 1” field of view would require d = 11 cm at0.55 µm. This would decrease the sensitivity of the instrument by about two magnitudes.Depending on what science topic to put emphasis on, field of view may be priviledged ratherthan sensitivity.An alternative technique would consist in mosaicing the field with patches of λ/d withdifferent pointings of the telescopes or differential FIRST tilts with the AO system loopclosed. This would work as long as the area of the object in a sub-field is bright enough toprovide 1 photon per exposure and per sub-pupil. This would work well on rather uniformextended objects like a galaxy but would not on a more compact source like a star + disksystem if one would like to image the disk.5.9.3 SensitivitySensitivity can be increased in different ways. As already done in this report, one way is toincrease the coherence of the wavefront using PUEO. This would improve the injectionefficiency in the fibers. The current PUEO would unfortunately not be suitable as most – ifnot all – of the visible light is picked-up for the wavefront sensor. The sensitivityimprovement would require to use a dichroic splitting the visible range, one part for PUEOand the remaining for FIRST. This requires to make a choice for the wavelength range forFIRST and the upper visible range, above 600 µm would probably be best given the sciencecases we have presented.Another technique is to increase the sub-pupil size but, because of the effect of saturation ofthe increase of injected light with sub-pupil size, this would just result in reducing the field ofview (see the section above). The increase in sensitivity would be effective only if PUEO wasused to increase the effective r 0 . But in this case the sensitivity would increase like the squareof the sub-pupil size. This would be well suited for faint point-like objects like AGNs forexample.5.9.4 Beam combinationOur baseline for beam combination does work. However, increasing the number of possiblecombinations may help improving the dynamic range. This needs further simulationshowever. One possibility is to consider integrated optics to concentrate a large quantity ofoptical functions in a small volume. We are investigating this possibility. It would probablybe for a second version of the instrument. An alternative solution for beam combination mayalso be studied to reduce the cost of the instrument as with the current baseline 7 detectors arerequired which drives the overall cost of FIRST. This is addressed in § 6.


495.10 Expected performance5.10.1 Summary of instrument and environment characteristicsThe instrument throughput is calculated in the spreadsheet below. Comments are selfexplanatory.We have studied the case with and without PUEO for AO correction. We haveassumed that part of the visible range would be used for PUEO but the other part would betransmitted without additional loss.Parameter Value Commentsseeing @ 0,55 µm (arcsec)0,65 Median valuer0 (m) @ 0,55 µm 0,17Effective r 0 with PUEO @ J (m)SR=27% on a bright object3,6Effective r 0 with PUEO @ V (m) 1,34PUEO OFF ON No throughput loss due to PUEOEffective r 0 @ 0.55 µm 0,17 1,34Instrument input:Subpupil diameter (m) 0,3Effective number of subpupils116 Ratio of areasµlenses: filling factor (hexagons-disks) 0,70µlenses: Strehl ratio0,95 Optical qualityAiry-to-fiber coupling efficiency0,82 No aberration caseµlenses: throughput 0,90Coherent energy 0,34 0,92 Fraction of turbulent psf coupled in fibers ( e "# $ 2 )Total throughput: 0,17 0,45Fibers:Fiber throughput1Fresnel loss per reflexion 0,04# reflexions 2Connector attenuation (dB) 0,3# connectors 2Total throughput: 0,80Material transmission!Instrumental visibility 0,87 0,87Single polarization effective throughput 0,50 0,50Cumulated polarizations effective throughput 0,71 0,71Inaccuracies of V-groovesPolarizations analyzed separatelyBoth are summed (photon noise-limited detection)Detector:Quantum efficiency 0,92Optical train# optics 10Total throughput:0,60Upper limit of number of optical pieces.95% efficiency per optical piece.Photometric throughput, unpolarized 0,07 0,20Coherent throughput, unpolarized 0,06 0,17Coherent throughput, single polarization 0,03 0,09Coherent throughput, cumulated polarizations 0,05 0,12


50We have not considered sky transparency in this analysis nor the telescope transmission. Butour extra train of optics should at least partly take that into account given that less than 10extra optics may be enough. This is to be confirmed or contradicted in the next step of thestudy.The instrument throughput is the input for the sensitivity and dynamic range studies. Up tofour cases are considered depending upon whether PUEO is used and whether polarizationsare used separately or are co-added after detection. From the above table one can read thatPUEO improves the injection efficiency in fibers by a factor 2.5.The instrument maximum snapshot field of view is given by the diffraction limit of a subpupilλ/d = 380 mas @ 0.55 µm.To be interferometrically achieved this field of view requires a minimum spectral resolutionof 12. If not the case for sensitivity reasons then the reconstructed field of view will besmaller than the maximum field of view.5.10.2 SensitivityLacour et al. (2007) have shown that images can be reconstructed as long as there is at leastone photo-electron detected per individual exposure. The sensitivity of FIRST depends onspectral resolution, wavelength and individual exposure time. We set the wavelength to0.55 µm and the exposure time to 4 ms (this latter parameter could be increased but we havemade the choice to be conservative and there is therefore some margin). The graphs below aretherefore indicative of the sensitivity for the parameters we have chosen.We have plotted two cases with R=12 (minimum resolution for maximum field of view) andR=250. As these figures show, the use of PUEO increases the sensitivity by 1 magnitude. Thecase where the two polarizations are cumulated has not been considered since images need tobe reconstructed in both polarizations first before they can be summed.The performances below could be improved using the multi-channel information toreconstruct wavefronts. The method is not yet developed and we have therefore consideredthat channels are analyzed independently.


51Given these hypotheses and choices, the sensivity of FIRST at 0.55 µm is:• V=11.7 with R=12 and V=12.7 if PUEO is used• V=8.4 with R=250 and V=9.4 if PUEO is used5.10.3 Dynamic rangePerrin et al. (2006) have shown that images are free of speckle noise and that the noise isuniform across the reconstructed image and consequently the dynamic range. The noise levelonly depends on photon noise. The dynamic range is thus increasing as the square root of thesource flux or as the square root of the integration time (the sum of individual imagesreconstructed for short exposures).We have performed simulations to calibrate this relation. We have simulated FIRST at CFHTand images of a binary system have been reconstructed. The contrast between the primary andthe secondary is 10 5 . The parameters used for the simulation are:• Wavelength: 0.55 µm• Bandwidth: 0.089 µm (R=6.2)• r 0 : 8 cm (seeing = 1,4”)• d: 26 cm (instead of 30 cm for the baseline solution)• V:5• Primary/secondary contrast: 10 5• Throughput: 100%Images below have been reconstructed with the algorithm of Lacour et al. (2007) and the Miraimage reconstruction software for integration times of 1s, 1min, 1h and 10h. A dynamic rangeof 2x10 6 (ratio of the bright peak to the rms residual) is estimated for the 10h exposure.Taking into account throughput and sub-pupil size for FIRST the following law can bederived:DR = 4.5 "10 6 ("10 #0.2" ( V +V loss #5)) $" int10 h " 6.2Rwith V loss the effective increase of magnitude induced by the losses in transmission of theinstrument (photometric transmission, polarization splitting, turbulence). V loss Depends on theuse or not of PUEO and on the addition or not of polarized images.!


52Results from simulation of the FIRST instrument. Image reconstruction was performed using the MIRAinterferometric image reconstruction software (Thiébaut, private communication). The object observed is amagnitude 5 star with a magnitude 17.5 companion assuming a 100% transmission for the instrument. Eachcolumn corresponds to different integration times. For the longest integration times, the dataset consist ofmillions of individual exposures. Reduction of the individual exposures was done recursively using the algorithmpresented in Lacour et al. (2007). The lower panels are diagonal cuts of the upper panels. The results show adynamic range increasing as the square root of the integration time. The companion is clearly detected after 1hour of integration time.The graphs below are computed for a total integration time of 5h and two spectral resolutions,R=12 and R=250.Whatever the case, the dynamic range is always larger than 10 4 beyond the limitingmagnitude in 5h integration time. For a 1 minute integration time (not presented here), it isalways larger than a few 100. For R=12 it is even larger than 10 5 for sources brighter thanV=8 in 5h of integration. This is not surprising as, by construction, FIRST is optimized forhigh dynamic range observations.


535.11 Compliance of baseline instrument with science specificationsThe performance of the instrument is summarized in the following table:Spatial resolutionField of view (snapshot mode with R≥12)Wavelength range (sorted by detectorquantum efficiency)23 mas @ 400 nm29 mas @ 500 nm34 mas @ 600 nm40 mas @ 700 nm46 mas @ 800 nm52 mas @ 900 nm275 mas @ 400 nm380 mas @ 550 nm619 mas @ 900 nm1. 500 – 700 nm, QE ≥ 90%2. 400 – 500 & 700 – 850, QE ≥ 50%3. 850 – 900 nm, QE ≥ 40%Spectral resolution 250channels can be co-added to reduce R at nocost (photon noise limited detection)R min =12 for maximum field of viewLimiting magnitude Without PUEO: V = 11.7 with R = 12V = 8.4 with R = 250Dynamic rangeWith PUEO: V = 12.7 with R = 12V = 9.4 with R = 250Uniform across the field≥ 10 4 in 5h integration timeas high as 7.10 6 at V = 0 in 5h≥ 10 3 for V ≤ 9 in 1 min integration timeas high as 7.10 4 at V = 0 in 1 min


54SpatialresolutionField ofviewWavelengthcoverageSpectralresolution /bandwidthResult of the compliance of the instrument with the science cases:Exoplanets Debris disks YoungStellarObjectsNot formally achievedbut detection possiblethanks to the highaccuracy knowledge ofthe PSF, the uniformityof the noise, and thesimplicity of the object(binary)Cepheids Novae NebulaeStellarSurfacesOK OK OK OK OK OK but edgyfor somephotospheres.No problemfor closeenvironment.OKOK for inner OK OK OK OK OK OKpart of disksOK OK OK OK OK OK OK OKOKOK forcontinuumNarrowfiltersrequired foratomic linesmV OK OK OKexceptforsourcesfainterthanV=13Dynamicrange# sources OK OK ~ OKTBCAGNOKOK OK OK OK OK OKOKEdgy forthe faintestphases butOKotherwiseOKOK with starsbrighter thanV=12.7 atminimumOK OK OK OK OK OK OK OKOKWork outwithV=12.7OKOK OK OKGenerally speaking the science case will be feasible with the instrument. Some restrictionsapply for some variable objects (novae) which may not be observable in a too-faint state orMiras which are bright but can also be very faint at their minimum.The lack of resolution for exoplanet detection will be compensated by the simplicity of theobject (point sources) and the excellent knowledge of the PSF, a characteristic of theinstrument. This super-resolution (well-known in interferometry, the extra resolution isderived from the a priori knowledge on the object) does not apply to the case of stellarsurfaces when below the resolution of CFHT because they are extended and unknown objects.The lack of field of view will be an issue for debris disks only to image their outer parts. Butit will not be to image the inner parts where planets form and where gaps are to be detected.5.12 Laboratory demonstratorIn this section, we describe the present status of the development of the laboratorydemonstrator. We have constructed a 36-element testbed in our lab which aims atdemonstrating the proposed technique. It employs a number of novel designs such as a


55segmented deformable mirror (DM) and silicon micro machined single-mode fiber arrays torealize precise star-light injection into fibers and beam recombination.Fig. 1: Schematic view of the experiment.5.12.1 OverviewThe experimental setup is shown in Fig. 1. To demonstrate an imaging capability of thisinstrument, an artificial binary star was used for this experiment. Two laser sources areinjected into 5 µm pinholes to make spatially coherent point sources and collimated to belaunched in the system. One source is used for an on-axis source (i. e. star) and another sourceis slightly tilted to simulate an off-axis source (i.e. companion). The two beams are thencombined by a beam splitter. A 37-element, segmented deformable mirror (DM) is placed inthe collimated beam to align the tip-tilt of each sub-aperture for precise injection into fibers.After reflection by the DM, the beam diameter is reduced a factor of 2.4 to match the pitch ofadjacent sub-apertures to the one of single-mode fibers. A microlens array divides the inputpupil to 36 sub-apertures and each beams is focused onto the fiber heads in the twodimensionalsingle-mode fiber array which consists of a bundle of 36 single-mode fibers.Output fibers are precisely aligned in a silicon v-groove chip and output beams from the fibersare recollimated by a linear microlens array. The beams are recombined in image plane tomake interference patterns. In the following sections, we describe the critical componentsused in this experiment.5.12.2 Injection opticsThe most important task for achieving very high dynamic range with this system is to realizean extremely high precision optical alignment of a bundle of 36 single-mode fibers. Foroptimal coupling, the position of the fiber core must be adjusted with sub-micron precision


56and we employed a combination of a two-dimensional fiber array, microlens array, and asegmented deformable mirror. A 2D single-mode fiber array consists in a bundle of 36 singlemodefibers fabricated by Fiber Guide Industries (Fig. 2). The fibers are packed in ahexagonal arrangement in a silicon substrate. The fiber pitch is 250 µm. The typical positionalerror is ~ 1 µm and the precision of fiber-to-fiber angularity is 2.5 mrad. The single-modefibers used for the 2D fiber array are manufactured by Oz optics. The fiber is a polarizationmaintaining fiber (PM-640HP) optimized for operation in the R-band with the followingparameters: diameter of the fiber core /cladding is 4 and 125 µm respectively; cutoffwavelength is 550 nm; numerical aperture is 0.12. The alignment accuracy of the slow axis ofthe polarization state in the fibers is ±3 degrees. However the precision of the fiber alignmentis not enough to achieve very high coupling efficiency and very high-dynamic range. We usea 37-element segmented deformable mirror based on microelectromechanical systemstechnology (IRIS AO) for very precise beam alignment with respect to fiber cores. Eachsegment consists of a hexagonal mirror attached to three individual electrostatic actuators,which provide tip-tilt and piston motion through differential actuation. An actuator has a 5 µmstroke. The largest inscribed aperture is 3.5 mm and the center-to-center distance between twoadjacent segments is 606.2 µm, including a 4 µm gap. This DM allows to compensate fiberposition errors in the 2D fiber array. At the moment, the DM is used for static positional errorcompensation, but it may be possible to correct the tip-tilt and piston arising from atmosphericturbulence.5.12.3 Beam combinerIn order to measure fringe visibilities, one needs to recombine the output beams from thefibers. We employed an image-plane combination scheme for this instrument. The fibers arearranged non-redundantly and one-dimensionally on a silicon V-groove chip from OZ optics(see Fig. 3), which is widely used for telecommunication. The uniform spacing of grooves(250±0.5 µm) gives precise alignment between fibers. The beams from the fibers arerecollimated by a linear microlens array whose lens pitch is 250±0.5 µm (48 element lineararray, SUSS MicroOptics), then spectrally dispersed by a prism. This will reduce the effect ofOptical Path Difference (OPD) and dispersion between fibers and furthermore it allowsspectro-interferometric observations. The beams are re-combined in the image plane tomeasure the object complex fringe visibilities.


57The non-redundant configuration is important so that any fiber pair has unique spacing.Unfortunately, due to a limited number of detector pixels, a non-redundant, 36-element onedimensionalarray is difficult to realize. To achieve sufficient fringe sampling, at least 4 pixelsare needed for the highest frequency fringes to avoid an aliasing problem and a visibility loss.In the case of a linear 36-beam combiner, the required number of linear pixels is over 10000.Therefore we divided the 36 fibers into 4 groups and each group uses one v-groove chip. Itmeans that 4 sets of 9-beam combiners are used for this instrument. For a 9-beam combiner, a48 channel v-groove chip is used and the fibers are placed on the following positions toensure non-redundancy: 2,3,7,14,27,29,37,43,46 (Fig. 4). The number of fringes inside a firstdark ring in the PSF is about 110, which leads to 440 pixels for sufficient fringe sampling.The 4 v-groove chips may share the same set of recombination optics.Fig.4. Configuration of the input redundant array and the output non-redundant array used forthis experiment. The black circles show the positions of the fibers.5.12.4 Spectrometer and Anamorphic opticsThe recombined beam is spectrally dispersed to minimize the effect of Optical PathDifference (OPD) between the fibers. Because the beams are recombined one-dimensionally,along the dispersion direction, the diameter of the Airy pattern corresponding to one spectralresolution must be small enough to avoid losing spectral resolution, while keeping the size inthe fringe direction. For a 9-beam combiner, the optimum diameter ratio (anamorphic ratio) is440:1.Our design of an anamorphic system is relatively simple (Fig. 5). It uses an afocalcombinations of two cylindrical lenses with the ratio of their focal lengths corresponding to amagnification of a pupil diameter, which translates into a small PSF in only spectraldimension in the image plane. On the other hand, another pair of cylindrical lenses having afocusing power perpendicular to the other two lenses is inserted into the beam, which leads toa large PSF in the fringe direction. Zemax simulations taking into account physical opticspropagation showed the maximum anamorphic ratio of this system is about 100:1 as shown inFig. 6. Strong spherical and chromatic aberrations of the microlens prevent one from reachingthe optimum anamorphic ratio.


585.12.5 DetectorThe detector integration time must be chosen so that atmospheric piston is frozen during anacquisition to keep the fringe contrast high. The typical atmospheric coherence time at visiblewavelengths is a few milliseconds. Therefore a key requirement for a detector is very lowread-out noise at very high frame rates for precise fringe visibility measurements. EMCCD(e.g. Mackay et al. SPIE 2001, 4306A, 289) is a cost-effective solution for this purpose. OurEMCCD camera, Hamamatsu C9100-13, can be read at high-frame rates (~ 31.9 frames/sec)for reading out a full area (512×512 pixel) and effectively very small read-out noise (


59simplicity. The configuration of the input redundant pupil and the output non-redundant pupilis shown in Fig. 5. Each input beam from a sub-aperture was precisely aligned with respect toa fiber for optimum fiber coupling by using the DM. Once the beams were aligned, the DMkept the mirror positions during the measurements. We measured fringes from the artificialbinary star at the He-Ne laser wavelength. The contrast ratio between the two source wasabout 10. Slow phase fluctuations were observed, probably thermally induced in single-modefibers. 100 frames of image data were collected and the integration time of each frame was0.1 second. We fitted a binary star model to the measured visibilities. Fig.7 shows measuredaverage fringe visibilities, closure phases and predictions from the best fit models. Visibilityand closure phase measurement accuracies were 2 % and 0.2 degrees respectively, which arevery good despite the existence of the phase fluctuations. The best-fit model is in goodagreement with the measured data, with a χ 2 /dof of 18 (dof=47). A contrast ratio of 15 derivedfrom the fitting well matched the directly measured value with a monitor CCD. Theseparation between the primary and the secondary source was 0.56 λ/d, where d is a subaperturediameter.In addition to the direct model fitting to the data, we reconstructed the original image by usingthe image reconstruction software MIRA (Multi-Aperture Image Reconstruction algorithm,Thiebaut, SPIE 2008, 7013, 7013-53). Fig.8 Shows the reconstructed image (right) and thedirectly measured CCD image (left). Two point-like sources are clearly visible in this image.5.12.7 ConclusionOur experiments have demonstrated the good data quality and the image reconstructioncapability of a pupil remapping system using a single-mode fiber for a single telescope.Thanks to the segmented deformable mirror, each input beam from a sub-aperture can bealigned very precisely with respect to a fiber and it allows excellent fiber coupling. Thelaboratory obtained fringe visibilities from the artificial binary star well matched thetheoretical model. Closure phases have been measured with 0.2 degree stability and fringevisibilities with 2% accuracy. Moreover, we successfully reconstructed the original imagefrom the visibilities by using the MIRA image reconstruction algorithm. This experiment arean important step toward realizing 36 fiber system on sky. We will attach this instrument to a1-meter telescope in Paris observatory next year.


Fig.8 Left: Directly measured CCD image. Right: Reconstructed image from the visibilities.60


616 Development requirements for the baseline projectThe concept of the instrument has been studied during the feasibility study. A few designshave been made and implemented in the laboratory prototype. But Further design is requiredfor the optics, the mechanics and the electronics of FIRST.6.1 Optical designWhat is missing as of today is a design for the pupil imager that will conjugate the CFHT3.6 m mirror and the segmented mirror of FIRST. In case FIRST is used behind PUEO, onewould need to re-image the pupil plane of PUEO.Also, no design has been made so far for the calibration source. This one is not very difficultbut needs to be done. In addition to its calibration purpose, the source will be both used as atest source in France and after shipping FIRST before it is mounted on the telescope. A basicspecification of the source is that it can remotely be set in and out of the beam.The optical design for the input of FIRST may have to be adapted to the future pupil imagerdesign.6.2 Thermal and Mechanical designFIRST will be set-up at the Cassegrain focus most likely downstream PUEO.One critical aspect for FIRST is the stability:• Mechanical stability to keep the instrument aligned: a 1 µm stability is required tokeep excellent injection in the fibers in the input and excellent mixing of the beams inthe output;• Thermal stability to avoid differential stretching or shrinking of the fibers which willcause differential dispersion.Our preferred solution is to have FIRST in a bonnette as a closed instrument to avoid dust.Heating will have to be reduced to a minimum to avoid internal turbulence and to ensurethermal stability. The detectors are the main heat source. A global budget needs to be workedout to take into account motors (steering mirror, segmented mirror) and additional heatingsources.The goal is to have the bonnette weigh a maximum of about 200 kg.6.3 ElectronicsMost electronics will be bought off-the-shelf. However, overall EMC (Electro-MagneticCompatibility) needs to be studied between the different subsystems. As well as to check thatthe experiment does not impact the telescope and his instruments. The electronic interface ofFIRST will comprise a check out of the global impedance and of the non-existence of groundloops.6.4 Data transfer and storageThe FIRST data will have to be transferred through the CFHT network. The experiment willproduce a high amount of data (see below). It will have its own data acquisition network(camera readings). Via an internal buffering stage, FIRST will post data to the local network.At the other end, both local telescope computers (for data tracking) and observers laptop


62computers will read and store data.Size of a window reading : 256 x 512 pixels.One window = one V-groove assuming all 2 (polarizations) x256 pixels are used in thespectral direction.Number of V-grooves = 14.Total number of windows: 14Exposure time : 4msDetector reading time : 30msAll windows are read every 34 msEach pixel is coded using 16 bits.A single read therefore amounts to: 14 x 256 x 512 x 16 bits ~ 30 MbitsAnd the data flow per second is: 0.86 GbitsFor a 1h continuous integration, the data flow is: 3 Tbits/hFor a full night, 30 Tbits will be produced.This is assuming the highest efficiency for the instrument and maximum spectral resolutionand for a continuous observation. However, sources requiring the longest integration timeswill be the faintest ones for which spectral channels will be summed and about 10 effectivechannels will be kept thus reducing the amount of data by 50. Also, the exposure time of 4 msmay be increased to improve the instrument sensitivity. 4 ms is the lower limit correspondingto fast turbulence.The data flow will clearly be part of the next study phase.6.5 Beam combinationThis is not necessarily a requirement in as much as the baseline solution works. However,the number of detectors required to detect the fringes is large (7) and the induced cost isimportant. In case a trade-off is to be made between science goals and cost, one could reducethe spectral resolution to a minimum and replace the spectrograph by multi-chromaticchannels. This would of course remove a lot of flexibility to the system like the effectivechoice of spectral resolution between 250 and just a few by summing channels as thedetection is photon noise limited. A solution is available to recombine 36 fibers in a 2Darrangement (Lacour, PhD thesis, 2007):78x78 pixels are required to sample the full fringe pattern. 4 such arrangements would berequired. Finding a solution to combine 42 fibers is certainly possible and would require about100x100 pixels. Only three of these would be necessary. An optical design would need to be


63studied to obtain chromatic channels across the 400 – 900 nm range each 33 to 75 nm wide toensure a maximum field of view. About 10 such channels would be needed and would fit oneach 512x512 pixel detector for each combination. This would work but this is not ourpreferred solution.The required development would be on the optical solution for the multi-chromator. Such adevelopment has already been carried out for the VEGA instrument of the CHARA array(formerly the REGAIN instrument of the GI2T interferometer, Mourard et al. 2001). 16channels are available. In this case no image anamorphosis is required and all in all thethroughput of the instrument does not change.


647 Before operations7.1 Tests in MeudonThe instrument will be fully tested in Meudon prior to shipping. Tests will be performed onthe calibration source to check the optical performance of the instrument (dynamic range andsensitivity specifications). A structure will be built to test the immunity of FIRST to telescopepointing and measure the residual flexures. A report will be written about these tests andprovided to CFHT to decide on instrument acceptance before shipping.7.2 ShippingMost instrument parts will be secured inside the bonnette to avoid a full disassembling of theinstrument. Parts to be removed from the bonnette will be carefully marked to avoid losingtime in the re-integration phase.The bonnette and the parts, as well as the test equipments, will be stored in crates equippedwith a shock tracking system. Shipping will be performed by a professional company and acustom broker will be chosen with CFHT to receive the crates in the best conditions possible.We plan a 2 month duration beetween the start of the disassembling of the instrument and thearrival of the crates at the summit.7.3 ReceivingOnce at the summit, everything is put back together with the help of notes and pictures takenin Meudon.Tests on the calibration source equivalent in Meudon will be performed in an integrationlaboratory to check the instrument status after shipping. Should parts be damaged duringshipping, this test will allow to identify and replace them. Same performance as measured inMeudon should be obtained at this point. Commissioning at the telescope will not start until itis the case.7.4 CommissioningAfter setting up the instrument at the Cassegrain focus, most likely attached to the PUEObonnette, tests can start. First of all, FIRST will be tested on its calibration source again tocheck that no problems happened during set-up. Then, it needs to be aligned with respect tothe CFHT optical axis both in image and pupil planes. A study of this procedure needs to beperformed in the next design phase as this will provide inputs about the required level ofremote control for the pupil imager. Once aligned, tests can be performed during daytime onthe PUEO internal source. These will be similar to the tests made in Meudon and in theintegration laboratory at the summit. Outcomes of these tests are:• Operability of FIRST at the telescope (remote control, data flow);• Performance on a sky calibrated source: first hint about the sensitivity of FIRST in realconditions;• Performance in the real environment: sensitivity to temperature, to telescope pointingdirection, to telescope mechanical environment, etc …• Handling of static aberrations at the telescope;• Instrument operation efficiency.


We anticipate that a single run of commissioning will be enough since the full instrument willbe tested in Meudon.65


668 Operations and data processing8.1 Set-upThe instrument will probably be set-up at the telescope only when it is needed forobservations and will be stored on the mezzanine otherwise. The procedure to set it up will besimilar to part of the commissioning tests. First, test on the calibration source. Second,alignment on the PUEO internal source. The first step may not be required if the instrument isstiff enough that a good level of alignment is kept between two consecutive runs.Since detectors do not require cryogenic fluids, the instrument can be quickly turn into activemode after shut down.8.2 CalibrationsPre-sky calibrations will consist in:• Pupil position measurement and centering on the segmented mirror of FIRST;• Optimization of individual segment tilts to compensate for static aberrations;• Measurement of static wavefront aberration and feedback to PUEO for compensation(equivalent to image sharpening);• Measurement of instrumental visibilities on the PUEO internal source (equivalent toPSF calibration);• Remote alignment of beam combiner if needed (3-axis stages to control V-grooves);• Spectrograph calibration on internal source.On-sky calibrations (no experience on this yet, will come after on-sky tests in Meudon later in2009):• Observation of point-like sources for PSF checks and data reduction algorithm checks;• The instrument is designed to be self-calibrated.8.3 Data acquisitionSee § 6.48.4 Data reductionA prototype software to reconstruct images has been made by Sylvestre Lacour and ÉricThiébaut and published in Lacour et al. (2007). It has been used to reconstruct images in thisreport. A data reduction pipeline will be delivered based on this prototype and on its upgrades.The language in which it will be written has to be agreed upon. The prototype is written inYorick, a free language used by part of the high angular resolution community.In addition, calibration reduction softwares will be added to the pipeline.8.5 Observing conditions requirementsFIRST will be very sensitive to atmospheric conditions. Computations have been performedassuming a seeing of 0.65”. Without PUEO, the instrument cannot be used with seeing poorerthan 1.5” as more than 8 magnitudes of sensitivity would be lost. With PUEO, one can expecta much smaller degradation of the performance. It is clear that the most demanding programswill require the median conditions at least as shown in this report.


679 Development roadmap and preliminary cost estimate9.1 Project scheduleUnless a rush is necessary, we plan to achieve first light on the 1m telescope in Meudonbefore we move on to a detailed design of FIRST for CFHT. As a matter of fact, we prefer togain some on-sky experience to avoid losing time when at CFHT. The current plan is to beready for sky tests in Meudon for the first semester of 2009. A detailed design of FIRST couldtherefore start in the second semester of 2009. The schedule would then be as follows:Detailed design phase: July 2009 – June 2010 (Final Design Review)Instrument construction: July 2010 – June 2011Tests in Meudon: July 2011 – September 2011 (report to CFHT for acceptance)Shipping: October 2011 – December 2012Integration at summit: January 2012First commissioning: March 2012 (leave one month in casemodifications or fixes are required)9.2 Manpower requirementsHere is a quite rough estimate of the manpower required to build FIRST per phase. This needsa deep study during the detailed design phase.Phase Topic FTEDetailed design phase Science case 1Optical design 0.25Mechanical design 0.5Thermal design 0.2Electronics design 0.5Simulations 0.25Data reduction 0.25Instrument construction phase Mechanics 0.5Optics (tests and alignment) 0.25Electronics 1Control software 1Data reduction software 1Tests in Meudon 0.5Integration at summit 0.2First commissioning 0.2Total manpower estimate 12


689.3 Cost estimateQuotations have been obtained for all off-the-shelf components (in red). Extra pieces ofmechanics, optics and electronics requiring a more detailed design have been estimated frompast experiences. This cost estimate needs of course to be consolidated in a next phase. Costsin Euros have been converted to US $ assuming 1 €=1.5 US$. All costs are rounded to thenearest upper 5 or 0. The following table is for the baseline solution (spectrograph).Part or subsystemQuantity Unit cost(k$)Total cost(k$)Segmented mirror and electronics 1 50 50Microlens array (input) 1 1 1Fiber bundle 1 25 25Fibers (dispersion matching) 126 0.1 13V-grooves 14 1 14Biprisms 7 1 7Filters 10 2 20Anamorphosis optics 7 4 283-axis stages for the V-grooves 14 2.2 31V-grooves µlenses 14 1 14HAMAMATSU C9100-13 7 39 273Atmospheric Dispersion Compensator 1 30 30Pupil imager 1 30 30Optical train (including mechanics) 1 20 20Beam steering mirror (electronics included) 1 20 20New dichroic for PUEO 1 10 10Internal calibration unit (point-like source, relay optics, 1 20 20motorized pick-up mirror)Mechanical interface with telescope (bonnette) 1 15 15Data acquisition and control PC 1 10 10Shipping costs 1 15 15Installation costs (3 weeks for 3 people) 1 6 18Commissioning costs (3 weeks for 3 people) 1 6 18Total estimated cost for baseline instrument: 682


69The multi-chromator solution yields the following budget (we have assumed that the cost ofoptical elements per 42-fiber beam combiner is the same as for 4 9-fiber V-groove completesystems including dispersive elements; as explained in the development requirement section,a more detailed investigation is needed as we may have overestimated the cost):Part or subsystemQuantity Unit cost(k$)Total cost(k$)Segmented mirror and electronics 1 50 50Microlens array (input) 1 1 1Fiber bundle 1 25 25Fibers (dispersion matching) 126 0.1 1342-fiber beam combiner 4 26,8 108HAMAMATSU C9100-13 4 39 156Atmospheric Dispersion Compensator 1 30 30Pupil imager 1 30 30Optical train (including mechanics) 1 20 20Beam steering mirror (electronics included) 1 20 20New dichroic for PUEO 1 10 10Internal calibration unit (point-like source, relay optics, 1 20 20motorized pick-up mirror)Mechanical interface with telescope (bonnette) 1 15 15Data acquisition and control PC 1 10 10Shipping costs 1 15 15Installation costs (2 weeks for 3 people) 1 6 18Commissioning costs (2 weeks for 3 people) 1 6 18Total estimated cost for multi-chromator instrument: 559


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Wyatt, M. C., Greaves, J. S., Dent, W. R. F., Coulson, I. M., 2005, “Submillimeter Images ofa Dusty Kuiper Belt around η Corvi”, The Astrophysical Journal, Volume 620, Issue 1, pp.492-500.76

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