`imaka: a path-finder ground-layer adaptive optics system for the University of Hawaii 2.2-meter telescope on Maunakea

Mark R. Chuna*, Olivier Laibc, Douglas Toomeyd, Jessica R. Lue, Christoph Baraneca, Simon Thibaultf, Denis Brousseauf, Hu Zhangg, Yutaka Hayanob, Shin Oyab

aInstitute for Astronomy, University of Hawaii-Manoa, Hilo, Hi 96720; bGemini Observatory, Hilo, HI 96720; cSubaru Telescope, National Astronomical Observatory of Japan, Hilo, HI 96720; dMauna Kea Infrared LLC, Hilo, HI 96720; eInstitute for Astronomy, University of Hawaii-Manoa, Honolulu, Hi 96822; fUniversite Laval, 2375, rue de la Terrasse, local 2104 Québec (Québec) G1V 0A6; gImmerVision, Montreal, Quebec H3A 2A5, Canada.

ABSTRACT

Astronomy with ground-layer adaptive optics systems will push observations with AO to much larger fields of view than previously achieved. Observations such as astrometry of stars in crowded stellar fields and deep searches for very distant star-forming galaxies pushes the systems to the widest possible fields of view. Optical turbulence profiles on Maunakea, Hawaii suggest that such a system could deliver corrected fields of view several tens of arcminutes in size at resolutions close to the free-atmosphere seeing. We present the status of a pathfinder wide field of view ground-layer adaptive optics system on the UH2.2m telescope that will demonstrate key cases and serve as a test bed for systems on larger telescopes and for systems with even larger fields of view.

Keywords: ground-layer adaptive optics

1. INTRODUCTION Adaptive optics (AO) systems are now commonplace at large ground-based observatories and they have enabled new observations across a range of astronomical fields. A new generation of optical and infrared AO concepts is emerging bringing with them the potential to broaden its impact. Ground-layer adaptive optics[1][2] (GLAO) is one of these concepts and it has the potential to dramatically increase the efficiency and capabilities of existing ground-based telescopes over a broad range of science observations by improving the seeing at the telescope sites[3]. By correcting only the optical turbulence close to the telescope (e.g. within the enclosure and in boundary layer), a GLAO system removes a significant fraction of the optical aberrations while providing a corrected field of view much larger than that of a classical AO system[4]. Recent measurements on Maunakea, Hawaii show that the wavefront phase variance is correlated over fields of at least half a degree[5][6][7]. As the next step to push GLAO to these fields, we are developing a “wide-field” GLAO demonstrator called `imaka with an areal field of view an order of magnitude larger than current or near-term GLAO systems. With it we will quantify the science gains and the performance of GLAO on fields of view approaching half a degree.

For astronomical science the key GLAO metrics are resolution, sensitivity, and field size. While many observations with current AO systems are focused on single objects, GLAO, with its wide fields of view, allow us to study thousands of objects simultaneously as well as extended objects subtending much larger angles. The science figures of merit typically used for AO (e.g. Strehl, wavefront variance, or contrast) are not as useful. A figure of merit better matched to the observations with GLAO is the rate at which signal-to-noise is gathered over the field of view. In particular we use the typical etendue, used in wide-field imaging surveys, divided by the image quality squared (A!/FWHM2) or better FOM=A!/NEA where A is the telescope collecting area, ! is the areal field of view, and NEA is the noise equivalent area[8]. The gains from GLAO with a very large field-of-view are now apparent. For a given site a 10-arcmin GLAO system on a 2-m telescope can collect SNR faster than a 2-arcmin GLAO system on an 8-m telescope. The potential impact for small telescopes is also apparent – a 4-m telescope with a 10-arcmin GLAO system and FWHM~0.3" is as capable as a 10-arcmin 8-m telescope with a seeing-limited FWHM~0.6".

In addition to having a different figure of merit than a traditional AO system, GLAO’s implementation is also different. With such large fields, some effects normally ignored in a classical AO system play defining roles in GLAO. For example, for a one-degree field of view visible-wavelength GLAO system, the range in altitude over which the turbulence is well corrected is small ~ 50 meters[2]; the deformable mirror (DM) altitude conjugation is critical. Similarly the combination of large field size and large pupil demagnification can result in a varying altitude conjugation

*mchun@ifa.hawaii.edu; phone 1 808 932-2317; fax 1 808 933-0737; ifa.hawaii.com

across the face of a DM tilted with respect to the chief ray in an optical relay. These error terms have not yet been quantified in a GLAO system on the sky. Also, contrary to intuition, GLAO is not a low-order AO system. To maximize the sensitivity, GLAO should correct for all the local turbulence. Its required correction order is set by the shortest science wavelength of interest. A GLAO system at V-band requires a ~ 20x20 actuator system on a 4-m class telescope and a kilo-actuator system on an 8-meter telescope. These technical challenges along with the push to wide fields of view provide incentive to develop visible-wavelength wide-field GLAO on smaller telescopes as a path to future systems.

2. SCIENCE CASE Our objectives are to demonstrate the technical performance and the scientific gains of wide-field GLAO. It is important to keep in mind that the number of fields we plan to observe is limited by our use of natural guide stars and by the observational efficiency of the system (or rather the lack thereof). When we observe with the instrument we will only have one or two target setups per night so deep integrations or observations of the same target field under a variety of observation/seeing conditions will naturally result.

Our first objective is to develop a GLAO system to demonstrate the performance of wide-field of view GLAO and measure how its corrected field, resolution, and sensitivity are linked to the phase aberrations arising in the atmosphere, dome, and telescope. This optical turbulence is not well characterized (e.g. turbulence origin, spectrum, and distribution) and is likely not simply replicated in the lab or in simulations. In quantifying GLAO performance we will address three key issues. First, we will quantify the gains of a GLAO system over very large fields of view in the visible and near-infrared. While previous studies have confirmed GLAO over modest fields of view (e.g. ~2-3 arcminutes) at near-infrared wavelengths[9][10][11][12], these field angles are still more than a magnitude smaller than what is predicted. With `imaka we will characterize GLAO over fields of view matched to and larger than the fields of view available on larger 8-m telescopes such as Subaru. More importantly, we will place this performance in the context of real, on-sky observing conditions.

Our second, overarching objective is to demonstrate science with GLAO over large fields of view by using the instrument for a selected set of observing programs. GLAO’s wide field of view and high spatial resolution will enable more efficient surveys, more detailed studies of crowded star fields, and more precise photometry and astrometry. We note that the astrometric power of the instrument comes from the combination of the gain in resolution/sensitivity and the ability to access more background astrometric reference stars. We have chosen the science cases below to highlight these aspects of GLAO and to quantify the system performance in metrics directly related to the scientific questions of interest. While the sky coverage for this system is limited, we have identified fields to demonstrate each of the cases.

Star Clusters in the Galactic Plane: Recent large-area near-infrared surveys, such as UKIDSS and VISTA, have resulted in many candidate young star clusters being detected within our Galaxy. Studying such clusters is important as most of our knowledge of star and cluster formation comes from a very small sample of star clusters within the limited environment (<500 pc) around the Sun. Some results suggest that there are two types of clusters formed: (1) dense, bound clusters and (2) “leaky”, unbound associations[13]. However, to determine whether these two types of clusters result from differences in primordial conditions or are different phases of an evolutionary sequence (in number of dynamical crossing times) requires an improved sample of clusters with well determined masses, ages, and radii in a range of environments[14]. While candidate clusters can be identified with seeing-limited surveys, given their distances (1-10 kpc) and the high stellar densities in the Galactic plane, their detailed characterization (mass, size, age) requires improved spatial resolution. In principle, characterization could be done with HST WFC3; however, the time available to mosaic the area covered by an individual cluster (>5') and observe many such clusters in different environments is prohibitive for HST.

Brown Dwarf Astrometry: Accurate measurements of proper motions and parallaxes of field brown dwarfs provide valuable insights into their physical properties. Proper motions combined with photometric or parallactic distances provide tangential velocities, which can be connected to the brown dwarf ages, either statistically for well-sampled groups of objects or even individually for extreme cases of high-velocity (i.e. halo) objects. Parallaxes are even more fundamental, as they establish the bolometric luminosities of brown dwarfs. Such astrometry requires high quality imaging over several arcminute FOVs, in order to have enough background reference stars for an astrometric reference frame. This cannot generally be done with standard AO (too small FOV) but is well matched to the astrometric gains provided by GLAO. We will use our GLAO system to measure proper motions and parallaxes of brown dwarfs in the

galactic plane. The resulting kinematics and distances will allow us to compare the properties with the known field population, e.g. their relative ages.

Exoplanet Masses: When combined with conventional radial velocity measurements, precision astrometric measurements will provide a direct measure of an exoplanet’s mass. We have selected several bright stars that are both known to have exoplanets based on the line-of-sight radial velocities and have the largest predicted astrometric wobbles; however, the astrometric signals are still extremely small, typically less than 500 micro-arcseconds, and can only be detected for stars that are nearby. Astrometric precision increases as the square root of the number of background stars, and inversely with the angular resolution; thus, we require the largest possible AO-corrected field of view to image a sufficient number of bright background stars. While exoplanet hosts are found all over the sky, there are not enough bright astrometric reference stars when the field of view is too small and there are NO nearby exoplanet hosting stars that have sufficient astrometric reference stars until we reach the fields of view obtainable with `imaka.

Dynamics of Dwarf Galaxies: Nearby dwarf spheroidals, such as Segue 1, give us a glimpse into the earliest epochs of galaxy formation based on their unusual abundance patterns, star formation histories, and masses[15]. Knowledge of the galaxies orbiting around the Milky Way is necessary to understand how the infall history of the galaxy may have influenced the galaxies evolution or our measurements of its properties. For instance, the measured internal velocity dispersions and resulting mass estimates may be overestimated due to past tidal encounters inflating the galaxy[16][17]. Furthermore, these tidal forces may have influenced the star formation history as the galaxy passed close to the Milky Way. Proper motion measurements of the individual stars within these galaxies can be used to determine the orbit of the galaxy and test how the Milky Way has influenced the evolution of low-luminosity dwarf spheroidals. The stars in these galaxies are distributed over >6', are too faint for GAIA (i~21), and should show proper motions 0.5 - 1.0 mas/yr[18]. Astrometric precisions per epoch should be ~0.5 mas. By combining proper motion measurements for many stars in a dwarf spheroidal galaxy, a further sqrt(N_stars) improvement is possible. Therefore, one can expect to achieve proper motion errors of 0.1 mas/yr for Segue 1, sufficient for measuring its orbit around the Milky Way.

Deep extragalactic surveys: With its excellent angular resolution over a very wide field our instrument will provide a unique tool for studying the evolution of galaxies and active galactic nuclei at a cosmological epoch, where both star formation and black hole growth peak (z=2-3). A number of well chosen fields with broad band multi-wavelength ancillary data already available (e.g. GOODS, COSMOS), in suitable locations for natural guide star GLAO can be observed for a very long time in several filters to study the morphology of distant galaxies and active galactic nuclei. In particular the sensitivity to faint diffuse features can be exploited to study the dynamic state and merger histories of these galaxies.

3. IMAKA CONCEPTUAL DESIGN In addition to the science goals outlined above, there are strong budgetary constraints that heavily impact the system implementation. As this is a pathfinder system we sacrificed operational efficiency and sky coverage in favor of a very simplified system that meets the basic science performance requirements. In particular we will observe using natural guide stars for all the wavefront sensors (impacts sky coverage), use existing hardware (sets order of the correction and sensitivity of the wavefront sensors), and will not include an atmospheric dispersion corrector (limiting observing zenith angles, wavelengths, and/or the bandpasses of the science filters. Even with these simplifications we meet the basic science requirements for a handful of fields on which we can demonstrate the science cases and the technical performance of wide-field GLAO. We leverage in-hand hardware and software to reduce the capital costs. We use the uncorrected f/10 Cassegrain focus of the UH2.2m Ritchey-Chrétien telescope. An optical relay reimages the telescope pupil to a deformable mirror and then the corrected beam is focused to the science camera. The baseline deformable mirror is a CILAS bimorph mirror and drive electronics from the now decommissioned Subaru AO36 system[19]. This mirror has 18 curvature actuators within the telescope pupil in two rings of 6- and 12-actuators and 18 “edge” actuators just outside the 60mm diameter design pupil for the mirror. The wavefront sensors are based on the wavefront sensors used in the mWFS experiment[7]. The WFS cameras consist of five commercial-grade EMCCD cameras from Raptor Photonics[20] and a set of off-the-shelf optics that form the pupil on the LLA and reimage the 8x8 Shack-Hartmann spot pattern to the detector. Each WFS subaperture is 4.5” in width with 10 pixels across each subaperture. Each WFS also contains a pupil imaging mode that images the pupil plane to the WFS camera to aid in system alignment. The mWFS computer will become the real-time control computer using a modified version of the RoboAO[21] AO control software. At the science focal plane, we

baseline a new large-area (10560x10560 pixel) CCD device from STA[22] being purchased for other purposes by the IfA. In addition, Don Hall (IfA/UH) is developing with Teledyne the next generation 4k x 4k 15-micron pixel HgCdTe devices (H4RG-15)[23]. IfA will have one of the first cameras based around an H4RG-15 and its development timescale is similar to `imaka. The main concession we made in consideration of the cost was the size of the field of view. During our past efforts, the optical design and implementation was identified as one of the major challenges for a GLAO system approaching a one-degree field of view. For this pathfinder we sought to minimize this risk but still implement a field of view that pushed past the field of any current (or near-term) GLAO system. There are two different drivers for the field of view. First the expected performance and the available science cameras set (1) the sampling and (2) the largest science field of view we can obtain in a single exposure. This is largely a practical limitation on the cost to fill a large focal plane. In the case of the H4RG-15 device and an expected FWHM in the near-infrared of 0.2-0.3”, we require a plate scale of about 0.1”/15-micron pixel. At this plate scale, the H4RG covers 7’x7’ and the STA10k device covers about 12’ x 12’. Second, since we are using natural guide stars for the wavefront sensing, we require a field large enough to acquire the guide stars and demonstrate the science and technical cases. For a field of view of 14’ diameter all of the science cases except the exoplanet astrometric case are possible. However, one of our goals is to demonstrate that much wider fields are accessible with GLAO so we have a strong desire to push the WFS guide star acquisition field to larger sizes (e.g. measure the GLAO correction on as large a field as possible. In addition, the optical design must deliver this field with a good pupil (distortion less than 1/10th of an actuator pitch for all fields) and image quality consistent with the expected GLAO correction over the science fields of view.

Figure 1- Two optical design concepts for imaka on the UH88" telescope. (LEFT) A refractive design with a field of view of 14' diameter and (RIGHT) a reflective design with a field of view of 24' diameter. Note that in the figure on the left the refractive design is shown with the UH88” wide-field corrector optics but this has since been removed from the design.

We developed two optical designs that met our basic requirements (see Figure 1 above). Our first optical design is based on a refractive 14’-diameter field which relays the telescope pupil to a 55mm intermediate plane and had diffraction-limited image quality over the center 7’ x 7’ field of view. This design meets the basic requirements, is cost effective (uses spheres and standard glasses in available sizes), and packages into a compact and stiff mechanical structure. Its ultimate field of view is limited by the size of the lenses we can readily obtain and the 14-arcmin diameter is a practical limit. Our second design is a reflective design that is based on an Offner relay but with the incoming and outgoing focal lengths being unequal (a “broken” Offner)[24]. With such a system the pupil moves off the Offner secondary mirror and into a position where the deformable mirror can be placed. All powered elements are spheres with the first element (AOM1) imaging the primary mirror to the deformable mirror (57mm diameter), the second powered element (AOM2) being a convex sphere, and the final powered element (AOM3) reimaging the field to f/13.3. The telescope f/10 to f/13.3 output focal ratio is well matched to the desired image quality and mechanical packaging constraints. The design was optimized for the image quality (RMS spot size) within the STA10k science field of view

(12’x12’) and provides an RMS spot size of about 12microns. This will limit the delivered image quality of the system under the best seeing conditions but we note that some tuning of the image quality is possible with the static figure on the DM as the optimization trades the field of view and the image quality at the center. Under the best seeing conditions a static figure correction on the DM of 1/7th micron coma and 1/3rd micron trefoil (phase) brings the image quality of the reflective design to be equivalent to the refractive design within the H4RG field of view. The main feature of the reflective design is that it provides a much larger WFS guide star acquisition field compared to the refractive design. The field is limited to about half a degree diameter by the off-axis aberrations the WFSs will see. The volume constraints at the Cassegrain focus of the UH88” telescope limit the largest mirrors (AOM3 and AOM4) to about ~16” diameter which provides about a 0.4 x 0.3 degree field. The wavefront sensors are quite tolerant since the imaging is made over individual subapertures and the dynamic range of each subaperture is large. The alignment tolerances of both designs are close to standard machine tolerances but on the reflective design the tilt of the final powered element is tight (1 arcminute). We down-selected to the reflective design driven principally by the desire to demonstrate GLAO on very wide fields of view and the addition of one science case (exoplanet masses) to the list of fields we can target with the instrument.

The complexity of the system is largely now in the mechanical structure and this is made worse by a weight restriction of 500 pounds (230kg) at the Cassegrain focus of the telescope. Our working concept is to sandwich the optics between two carbon-fiber plates with the mirrors mounted to walls and bulkheads between these two side plates. The main challenges we see in this approach are (1) the initial positioning/alignment of the optics and (2) and the dimensional stability of the instrument. Unlike a machined structure, we do not believe that the structure can be fabricated with interfaces positioned to meet the tolerances of the mirrors. We instead plan to measure the mechanical interfaces of the structure with respect to a common reference point, and then adjust the interface between the structure and the mirror mounts to position the mirrors as needed. The overall positional tolerances of each of the mirrors is roughly 100 microns and we believe the structure will be built with interfaces accurate to about 3x this. This work will be done during the integration at sea level (e.g. Hilo, Hawaii). The structure will be designed to be rigid with respect to the changing orientation when it is at the back of the telescope and in this respect we believe the CF composite structure will surpass the performance of a similar structure made from aluminum. This is true also for the dimensional changes of the system due to temperature changes since the coefficient of thermal expansion for carbon fiber composite is closer to that of glass than aluminum. However, carbon fiber composites are hygroscopic so exhibit dimensional changes due to changes of moisture content in the composite. The scale of the coefficient of moisture expansion (CME) effects are similar to those of the changes of aluminum due to thermal changes but the moisture expansion changes occur over much longer timescales (days to weeks) as opposed to thermal changes (minutes to hours)[25]. Over the course of a night the dimensional changes due to CME should be negligible. We believe this will be beneficial in terms of the calibration of the focal plane distortions (for astrometry) and of the image quality (PSF stability) on any given night. However, it remains to be demonstrated that the system calibrations will be sufficient to combine observations over very long periods (months to years). The largest changes will occur in the transitions between the integration facility (at sea level) and the summit where the change in water content in the air is large.

A calibration unit (CalUnit) at the front focus of the instrument (telescope focal plane) will be used to assess the image quality and the focal plane distortion and provide artificial guide stars for the AO alignment and calibrations (e.g. interaction matrices), the initial WFS centroid offsets in the individual WFSs, and the functional checks for the AO subsystems. The CalUnit will be available for use during the instrument integration and when the system is on the telescope. In addition, the science field calibrations and the simulated AO WFS guide stars can be deployed independently at any time. We expect this functionality to be useful when making astrometric calibration tests on sky with the AO system locked to a fixed DM shape (e.g. seeing limited calibrations). The baseline astrometric calibration, however, is to deploy a pinhole grid over the science field of view in the CalUnit. This will be dithered about the field to provide a distortion correction and calibrate and changes with environment (temperature, humidity) or observing condition (elevation angle). This approach is similar to the baseline approach for the TMT first-light instrument IRIS/NFIRAOS[26]. We note also that in our case the telescope focal plane is curved and contains off-axis aberrations that will not be implemented due to cost in the initial version of the imaka CalUnit. As a result of the flat pinhole grid, the focus error in the final focal plane will vary across the field. The amount of defocus is minor for the H4RG detector but noticeable over the STA10k field. In addition, since the telescope focal plane (entrance focal plane of the instrument) is not telecentric, the pinhole spacing will effectively change across the field (even if the pinhole grid is perfectly uniform). All of these effects will need to be carefully measured and calibrated as part of the instrument integration and commissioning.

4. PERFORMANCE ESTIMATES Performance simulations have continued to evolve from previous reports[27] but we reach the same basic conclusion that the free-atmosphere seeing is attainable over wide fields of view with GLAO at Maunakea. We have updated the optical turbulence profile using the Gemini MKGL study for the turbulence within the ground layer (first few hundred meters)[4] and an analysis of the MASS running at the summit of Maunakea from 2009 through 2013. The resulting ten-layer model was binned from the full MKGL ground-layer profiles and the MASS (12-layer model) with the integrated strength and the 5/3rds moment of the CN

2(h) profile preserved in the binning. The profile extends to about 15km to be useful in simulations of GLAO with laser guide stars (e.g. Rayleigh and sodium). We also updated the wind velocity profile based on the measured ground wind speeds at the CFHT/Gemini weather tower (10 years), measurements of the dome seeing at CFHT and the UH88”[7], and the Global Forecast System upper wind velocity predictions above Hawaii. The optical turbulence and velocity profiles are given in Table 1 for the median ground-layer and median free-atmosphere conditions. The integrated conditions at 0.5um for this profile are r0=0.128m, seeing=0.81”, and a "0 = 2.7”.

Table 1- imaka Mauna Kea median(GL)/median(FA) optical turbulence profile

Altitude (m above site)

Cn2dh (m1/3)

Fractional Strength Speed (m/s) Direction

(0-N, 90-E) 0 (dome) 16.8e-14 0.365 0.4 0 0 11.5e-14 0.250 3.5 90 15 4.51e-14 0.098 6.5 90 37 2.73e-14 0.059 6.5 90 197 1.02e-14 0.022 6.5 90 357 0.265e-14 0.006 6.5 90 804 2.08e-14 0.045 6.5 90 3171 1.87e-14 0.041 9.4 260 7307 3.02e-14 0.066 22.1 260 14636 2.20e-14 0.048 5.9 260

Limitations in the implementation of the GLAO system are important. For GLAO implementations on large telescopes, an adaptive secondary is a reasonable solution – the misconjugation of the secondary mirror is not problem for fields up to about 20 arcminutes[28]. For this pathfinder such a solution is not affordable and in addition we seek to push to even larger fields where the conjugation of the deformable mirror is important. The simulations now include the residual due to the GLAO correction including a model of the DM correction, the distortion of the pupil and misconjugation of the DM due to the tilt of the DM with respect to the pupil, the WFS sensing and wavefront reconstruction (average), centroid offsets of the WFSs due to off-axis aberrations in the optics, and also the static aberrations in the optical design that are not corrected by the DM. The expected FWHM and Strehl from the imaka Monte-Carlo simulations are shown in Figure 2 below. In the NIR the imaging performance is impressive with modest Strehl Ratios possible over very large fields. At these wavelengths the PSF from the residual atmosphere and the diffraction limited PSF are comparable and set a floor for the GLAO correction at longer wavelengths on this telescope.

Figure 2- Estimated imaka performance in FWHM (left) and Strehl (right) on the UH88" telescope as a function of field position and wavelength within the STA10k field of view under median seeing conditions by O. Lai. The wavelengths shown are g(480nm), r(622nm), i(760nm), z(905nm), Y(1020nm), J(1.2um), H(1.65um), and K(2.2um) (purple thru red).

The degradation of the image quality at the edge of the field is due to the optical design degrading in one direction at the edge of the science field of view. The GLAO “atmosphere only” performance is flat over this field of view.

The simulations presented above do not include time lag error, atmospheric dispersion, and some non-common path errors. For the first, we estimate the time lag error to be small due to the slow rate at which the ground-layer and dome seeing change[7]. For the second, for the GLAO performance tests we will use a medium band filter bandpass, observe close to the zenith, and/or work at longer wavelengths. In the NIR the atmospheric dispersion across the major bandpasses are completely negligible in comparison to the residual phase variance. The non-common path errors in a GLAO system include the field aberrations (which normally are negligible in an AO system). The design field aberrations are included in the simulations but the figure errors on the individual mirrors and the more normal non-common path errors (such as aberrations in the WFSs themselves) are not included. Since the WFSs pick off the natural guide stars very close to the final focal plane errors in the “science path” are negligible and our experience with the mWFS wavefront sensors are that the WFS static aberrations are very small. The major non-common path errors are therefore already included in the simulations.

The largest uncertainty in the performance estimates is the outerscale or spectrum of the optical turbulence in the dome. For the simulations to date we have used a wavefront outerscale of 30m but there is some evidence that at the UH88” this may be smaller[7]. If this is the case the performance at shorter wavelengths may degrade due to an increasing fitting error.

5. NEXT STEPS We are now entering the detailed design phase of a wide-field pathfinder GLAO system for the UH88” telescope. The detailed mechanical structure and mirror mounts are expected to be complete by the end of the summer and the rest of the system (controls, electronics, and calibration unit) by the end of the year. Most of the hardware is already in hand and the procurement is driven by the structure and the optics themselves. System integration in the lab will be in approximately in one year and we plan to start bringing the instrument to the telescope in the fall 2015.

Observations are planned in short campaigns where we will bring the instrument to the summit, conduct a set of observations/tests, and then bring the system back down to the lab at sea level for further improvements/modifications. The sequence of on-sky experiments and observations will progress from basic GLAO performance characterization to the more detailed science case characterization. We expect to observe with the instrument over several semesters to fully characterize its performance and its science gains.

Potential future upgrades include increasing the sky coverage of the instrument for more regular science use and pushing to shorter wavelengths. A higher order DM and an atmospheric dispersion corrector would improve the performance in the visible but to enable regular science observing we will need to increase the efficiency and sky coverage of the instrument. Our expectation is that we will deploy multiple UV Rayleigh guide stars in a scheme similar to that used in the RoboAO system[29]. The current optical design supports UV Rayleigh laser guide stars though the current Subaru AO36 DM has a silver coating so it would need to be recoated/replaced. Simulations show that low-altitude Rayleigh laser guide stars provide better GLAO performance than the same number of natural guide stars or sodium laser guide stars due to the faster averaging of the residual upper atmosphere in the GLAO reconstruction. In addition, if the low required temporal bandwidth of GLAO is confirmed then we could potentially work with a single laser that is dithered about the field.

ACKNOWLEDGMENTS

The current `imaka work is funded by the National Science Foundation (AST-1310706) and previous work on `IMAKA was funded by the Canada-France-Hawaii Telescope. The current work is intimately tied to the CFHT `IMAKA studies we are indebted to the CFHT team for their support and efforts.

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