The Gemini Planet Imager: First Light

The Gemini Planet Imager (GPI) is a dedicated facility for directly imaging and spectroscopically characterizing extrasolar planets. It combines a very high-order adaptive optics system, a diffraction-suppressing coronagraph, and an integral field spectrograph with low spectral resolution but high spatial resolution. Every aspect of GPI has been tuned for maximum sensitivity to faint planets near bright stars. During first light observations, we achieved an estimated H band Strehl ratio of 0.89 and a 5-sigma contrast of $10^6$ at 0.75 arcseconds and $10^5$ at 0.35 arcseconds. Observations of Beta Pictoris clearly detect the planet, Beta Pictoris b, in a single 60-second exposure with minimal post-processing. Beta Pictoris b is observed at a separation of $434 \pm 6$ milli-arcseconds and position angle $211.8 \pm 0.5$ deg. Fitting the Keplerian orbit of Beta Pic b using the new position together with previous astrometry gives a factor of three improvement in most parameters over previous solutions. The planet orbits at a semi-major axis of $9.0^{+0.8}_{-0.4}$ AU near the 3:2 resonance with the previously-known 6 AU asteroidal belt and is aligned with the inner warped disk. The observations give a 4% posterior probability of a transit of the planet in late 2017.


Significance Statement
Direct detection-spatially resolving the light of a planet from the light of its parent star-is an important technique for characterizing exoplanets. It allows observations of giant exoplanets planets in locations like those in our solar system, which are inaccessible to other methods. The Gemini Planet Imager (GPI) is a new instrument for the Gemini South telescope. Designed and optimized only for high-contrast imaging, it incorporates advanced adaptive optics, diffraction control, and a near-infrared spectrograph and an imaging polarimeter. During first light scientific observations in November 2013, GPI achieved contrast performance that is an order of magnitude better than conventional adaptive optics imagers. Adaptive Optics | extrasolar planets | astronomical instrumentation D irect imaging is a powerful complement to indirect exoplanet detection techniques. In direct imaging, the planet is spatially resolved from its star, allowing it to be independently studied. This capability opens up new regions of parameter space, including sensitivity to planets at > 5 AU. It also allows spectroscopic analysis of the light emitted or reflected by the planet to determine its composition [1] [2] and astrometry to determine the full Keplerian orbital elements [3] [4].
Imaging planets is extremely challenging-Jupiter is 10 9 times fainter than our sun in reflected visible light. Younger extrasolar planets are more favorable targets. During their formation, planets are heated by the release of gravitational potential energy. Depending on the exact formation process and initial conditions, a 4 Jupiter-mass (MJ ) planet at an age of 10 million years could have a luminosity between 10 −6 and 2 × 10 −5 L [5]-but this is still a formidable contrast ratio. To overcome this, astronomers have combined large telescopes (to reduce the impact of diffraction), adaptive optics (to correct for phase errors induced by atmospheric turbulence), and sophisticated image processing [6] [7]. This recipe in various combinations had achieved several notable successes [8] [9][10] [11] [12]. However, the rate of these discoveries remains low [13][14] [15], in part because the number of suitable young stars in the solar neighborhood is low, and, for all but the closest stars, detection is limited to > 20 AU where planets may be relatively rare. To move beyond this limited sample, dedicated instruments are needed that are designed specifically for high-contrast imaging. One such instrument is the Gemini Planet Imager (GPI). GPI is a fully-optimized high contrast AO facility deployed on an 8-m telescope and is almost an order of magnitude more sensitive than the instruments of the previous generation. With this powerful dedicated facility, large-scale surveys could increase the sample of directly imaged giant planets to 25-50, or more [16]. We present an overview of the design of GPI, discuss its initial operation and performance, and show first-light science results.

Design of the Gemini Planet Imager
At high dynamic ranges, the point spread function (PSF) of a bright star imaged through a telescope is complex [17][18], with light scattered by a variety of different mechanisms and imperfections interacting coherently and varying on different timescales. Atmospheric turbulence induces large wavefront phase errors (and smaller intensity fluctuations.) These are corrected (imperfectly) by the adaptive optics (AO) system; the residual phase errors produce a halo of scattered light. Instantaneously and monochromatically this halo is broken up into individual speckles. Over long exposures, these speckles move randomly and the atmospheric component of the halo reduces to a smooth profile. However, if the measurements of the phase used to control the adaptive optics are biased, the residual phase will have a non-zero mean and the speckle pattern will have persistent components that could be mistaken for a planet. These can be attenuated through PSF subtraction [6], but often the exact bias will vary as, e.g., the mean seeing changes, resulting in a halo that changes on timescales of minutes [19]. An additional source of persistent speckles are intensity variations either due to reflectivity variations or to propagation of wavefront errors on surfaces not conjugate to the deformable mirrors (DM) [20]. Even in the absence of wavefront errors, the finite aperture of the telescope causes diffraction that can swamp the planetary signal.
GPI is designed to deal with these errors through multiple approaches. The AO system has a large number of degrees of freedom and operates at high cadence to minimize residual atmospheric turbulence. An apodized-pupil Lyot coronagraph (APLC) [21] removes residual diffraction. Most importantly and uniquely, the entire design of the optical system is intended to minimize biases in the wavefront measurement and in turn minimize quasi-static speckles.
The instrument was also designed to minimize the chromaticity of the PSF-the APLC design sacrifices inner working angle for achromaticity [21], the number of transmissive elements is kept to a minimum, individual mirrors are polished to < 1 nm root mean square (RMS) surface error 1 , and optics near focal planes are avoided to reduce chromatic Fresnel effects. As a result, the residual speckle pattern is highly correlated between different wavelength channels and can be further attenuated through multi-wavelength postprocessing [22]. Table 1 summarizes the properties of the instrument.
Adaptive optics and optomechanical systems. The adaptive optics system is intended to improve on previous facilities in two respects-lower total wavefront error from dynamic sources, and also lower quasi-static errors by an order of magnitude. Typical current AO systems such as the 349-actuator system on the Keck telescopes [23] has 250 nm RMS of dynamic wavefront error on a bright star [34]; GPI is designed to achieve 90 nm. The static errors of current systems are hard to estimate, but generally are 50 to 100 nm RMS. For the Keck AO system this number is 113 nm [24]; the design goal for GPI was 10 nm.
To achieve this goal, GPI adopts several novel features. The high-order AO system has 43 subapertures across the the Gemini South pupil. A conventional piezoelectric DM with this number of subapertures would have a diameter ≥ 20 cm, resulting in an instrument too large for the Cassegrain focus. Instead, GPI uses a silicon MEMS DM manufactured by Boston Micromachines corporation [25]. This DM has 64 × 64 (4096) actuators-GPI uses a 48-actuator circle (including slaved actuator rings) within this area. This compact (25.6mm diameter) device enables precision AO correction within a small volume and mass (2200 kg). The MEMS DM has two noteworthy limitations. First, the total displacement range is ∼ 4 µm, insufficient to fully correct atmospheric turbulence. GPI therefore also employs a piezoelectric, low-order DM, referred to as a "woofer." Second, the MEMS DM has five defective actuators inside the controlled pupil. Light from these bad actuators is blocked by the coronagraphic Lyot mask. Tip/tilt control is provided jointly by a piezoelectric mount for the whole woofer DM (< 30 Hz) and actuation of the woofer surface itself (> 30 Hz).
GPI employs a Shack-Hartmann wavefront sensor (AOWFS). Such sensors are susceptible to aliasing of errors outside the Nyquist range, which can inject significant midfrequency wavefront errors into the final image. A spatial filter [26] was installed to mitigate this. The spatial filter is adjustable with size set by the Nyquist criterion. The spatial filter can also be used for calibration. By adjusting it to a pinhole we can inject a nearly perfect spherical wavefront into the AOWFS, allowing fast calibration of the dominant source of non-common-path aberrations in the system. GPI uses a computationally efficient Fourier transform reconstructor [27] to translate AOWFS slopes into DM commands. The control loop gain of each Fourier mode is individually optimized [28] to minimize total wavefront error as atmospheric turbulence and star brightness vary.
GPI operates at a Cassegrain focus and so is subject to a variable gravity vector. The mechanical structure of GPI is therefore designed to minimize and actively control flexure. Individual components such as the DM and lenslets remain optically registered to each other to within a few microns.
Coronagraph, science instrument, and IR wavefront sensor. GPI uses an APLC, which suppresses diffraction by using a grey-scale mask to taper the transmission near the edges of the pupil. The transmission profile of this mask is matched to the Fourier transform of the hard-edged focal-plane stop (radius 2.8λ/D) such that residual diffraction is channeled outside the telescope pupil at a Lyot stop [21]. The masks were implemented with a dithered half-tone pattern of 10 µm metal dots on a glass substrate; the focal-plane stop operates in reflection as a gold-coated mirror with a small central hole. The machined Lyot stops are located inside the cryogenic spectrograph. A grid of narrow, widely-spaced lines printed onto the apodizer forms a two-dimensional grating, producing diffracted images of the central star in a square pattern. These satellite images (Fig. 1) track the star's position and intensity, facilitating astrometric and photometric calibration data [29] [30]. A flexible and automated data pipeline transforms raw data into calibrated datacubes and starlightsubtracted images. Steps in this process include removal of detector systematics, wavelength calibration, compensation for time-variable flexure, assembly into multiwavelength datacubes, astrometric and photometric calibration into physical units, and PSF calibration and subtraction. This open-source software and extensive documentation is freely available.
A near-infrared (IR) pointing and low-order wavefront sensor maintains the star in the center of the focal plane mask hole through feedback to the AO system. GPI also includes a high-accuracy near-IR interferometer integrated with the coronagraph optics to provide time-averaged measurements of the wavefront at the focal plane stop, similar to that used with the P1640 instrument [31], though this has not yet been tested on sky.
The science instrument is a near-IR integral field spectrograph (IFS) [32]. Similar to the OSIRIS and P1640 instruments, it samples each spatial location in the focal plane using a lenslet array and disperses each resultant sample using a prism producing ∼ 37, 000 individual microspectra. Each spatial pixel corresponds to 0.0143 arc seconds on the sky. In a single exposure, for each location in the field of view the IFS produces a λ/δλ 30-70 spectrum over one of the standard astronomical bands (Y JHK, with K-band split in two.) The IFS uses a Teledyne Hawaii 2RG detector. Operating in upthe-ramp sampling mode it has a readout noise of 5 electrons RMS. It is cooled to ∼ 80 K with a pair of 15 W Stirling-cycle cryocoolers.

First-light results
GPI was extensively characterized at the University of California, Santa Cruz prior to shipping to Chile [33]. In October 2013 it was attached to the bottom Cassegrain port of the Gemini South telescope. After some brief alignment tests, first light occurred on November 12, 2013 (UT); the first testing run continued until November 18, with a second observing run from December 9 to 12.
AO performance. The instrument met its initial performance goals. The AO system locked onto every target attempted, even in seeing with a Fried parameter of r0 5 cm. The AO loop closed on stars I ≤ 8 mag. (GPI operates down to I 10 mag. in better seeing.) GPI can save extensive AO telemetry at the full system frame rate, including DM commands, raw wavefront sensor images and centroids, reconstructed phases, and tip/tilt measurements and commands. These have been used to generate an error budget ( Table 2) for typical observations of β Pictoris (I = 3.8 mag.) from observations taken on December 11 (UT). The Fried parameter r0 was estimated to be 20 cm from measurements earlier in the night with a Differential Image Motion Monitor located outside the Gemini dome, corresponding to better-than-average seeing. Most error terms are directly estimated from AO telemetry, following methods used in [34].
Two error terms are worth further discussion. The first is a quasi-static offset on the MEMS DM, consisting primarily of a phase excursion in the outer two rings of subapertures. Although this is the largest quasi-static term in the error budget, it likely has only a small effect on final images, since the APLC transmission for those subapertures is low. The likely cause is a reconstruction of centroid patterns that do not correspond to a physical wavefront (such as rotation of the whole centroid pattern), and it will be mitigated through software changes. The second is a strong 60 Hz vibration. The Stirling cycle cryocoolers on GPI induce vibration at harmonics of 60 Hz in the GPI mechanical structure. Internal to GPI, these contribute about 6-8 mas of RMS tip/tilt image motion at 120 and 180 Hz, as measured with our artificial star unit. Observing actual stars through the telescope, we see an additional 10 mas of 60 Hz image motion, indicating a component external to GPI is vibrating. More significantly, there is 100 nm (time-domain root-sum-squared of the spatial RMS) of low-order wavefront vibration at 60 Hz, primarily focus, when the cryocoolers are operating. This may indicate that the telescope's mirrors, which are very lightweight, are being driven by the GPI vibrations. We carried out measurements both with the cryocoolers at full power and reduced to 30% power. Modifications to the cryocoolers in January 2014 are expected to significantly reduce the vibration.
With the cryocoolers at full power, we estimate the RMS wavefront error to be 134 nm; with the cryocoolers at minimum power, the RMS wavefront error is 98 nm (exclusive of tip/tilt), corresponding to a Strehl ratio of 0.87 at 1.65 µm. Although this is very good AO performance, it is not vastly better than other high-order AO systems. However, GPI significantly exceeds the raw image contrast of such systems-we ascribe the difference to much better control of quasi-static and systematic errors.

First-light science: Beta Pictoris b
Here we discuss observations of β Pic. This young, nearby (19.4 pc), A6V star has a bright, edge-on debris disk [35] and a directly-imaged super-Jupiter, β Pic b, orbiting at 10 AU [10]. Various asymmetries in disk structure have been discovered [36] [37], including a midplane warp at ∼50 AU that has been attributed to a possible planetary perturbation [38][39] [40].
The planet has been detected by VLT/NACO [10], Gemini/NICI [41], and the Magellan AO system [42]. We observed the planet in H band (1.65 µm) on November 18 (2013) UT. We obtained 33 individual 60-second images in coronagraphic mode, with the cryocoolers operated at minimum power. The planet was immediately visible in a single raw 60-second exposure (Fig. 1.) For comparison, a lower signal-to-noise H-band detection using NICI [41] required 3962 s of exposure and extensive PSF subtraction.
A set of 30 images were PSF-subtracted using a modified version of the TLOCI algorithm [43]. For each data cube wavelength slice, the remaining wavelengths, and the images at the same wavelength acquired at a different time, are used as PSF reference images. A simulated planet having either a flat spectrum similar to L-dwarfs or a CH4-dominated spectrum similar to T-dwarfs is used to exclude reference images with more than 10% self-overlap of the planet with the image being processed. The selected reference images are then used to subtract the stellar noise by performing a LOCI least-squares fit constrained to positive coefficients. All starsubtracted data cubes are then rotated to put North up and median combined (Fig. 2.) We evaluate the contrast of the final image by calculating the standard deviation of the intensity in concentric annuli about the star. Figure 3 shows a plot of contrast for a 30minute set of GPI data. In 30-minute exposures, GPI is performing ∼ 2 magnitudes better than typical contrast obtained at the Keck Observatory.
Astrometry. We use these images to constrain the orbit of β Pic b. Most astrometric calibrators are unsuitable for GPI because of the small field of view (2.8 arc seconds on a side) and also because for low-contrast ratio binary stars the secondary will either saturate the science detector or light from the secondary will perturb the wavefront sensor. Therefore, we calibrated the scale and orientation of GPI using other extrasolar planets as references. We compared the offset between HR 8799c and HR 8799d measured at the W.M. Keck Observatory in October 2013 with GPI observations of this system. The position of β Pic b is measured from the combined TLOCI image. The position of β Pic b is estimated in each raw 60-s image from the center of the square pattern defined by the four reference-grid spots. The measured separation of β Pic b is 434±6 mas at a position angle of 211.8±0.5 degrees.
We fit the new GPI astrometric measurements presented here together with the compilation of data in [3]. We use Markov Chain techniques described in [4]. The adaptive Metropolis-Hasting sampler used in [4] has been replaced by an affine invariant sampler [44] [45]. Figure 4 shows an ensemble of trajectories computed using orbital elements sampled from the joint posterior distribution. The results show a good fit with a posterior expectation value of χ 2 ν = 0.76 (ν = 14). The GPI data point implies that β Pic b has recently turned around on its orbit and roughly one half an orbital period has elapsed since the first astrometric measurement in 2003. The combination of small measurement errors and the lever arm provided by extended orbital coverage yields the best estimate to date of the semi-major axis. Our new measurement places β Pic b at a semi-major axis of 9.0 +0.8 −0.4 AU, with a corresponding period of 20.5 +2.9 −1.4 yr. Beta Pic b will remain observable (> 0.2 arcseconds) at least until March 2016 and will likely reappear before October 2019. A possible transit of the planet across the star was observed in 1981 [46]. Our estimated orbital inclination is 90.7 ± 0.7 • ; transits of β Pic will occur if the inclination is within 0.05 • of edge on. The next opportunity to observe such a transit by β Pic b is in September through December 2017 (68% confidence).
The orbit has a small but nonzero eccentricity (e = 0.06 +0.07 −0.04 ). The corresponding range between periaspe and apoapse (8.8-9.4 AU; see Table 3) falls neatly between the 6.4 AU and 16 AU asteroidal belts [47][48] [49] and places the inner belt close to the 3:2 commensurability with the orbital period of β Pic b Our results agree with [50][3]-the chief difference is the improvement in the confidence intervals (see Table 3). Distinct peaks are beginning to emerge in the posterior distribution for epoch of periapse and for ω (which is significant for identifying the perturber of the in-falling evaporating bodies.) The 68% confidence interval for i drops from 3.4 • to 1.36 • . The posterior probability of transit has increased by a factor of 2.2 relative to [3], but the chance that β Pic b transits the star is still relatively small (4%).
The longitude of the ascending node now has a confidence interval of 0. • 9 compared to 3 • previously. The corresponding PA is 211.6 ± 0. • 5 and clearly misaligned with the main disk (PA = 209 ± 0. • 3) but more nearly aligned with the inner warped component [3] [51].
These results show the power of a dedicated high-contrast imager; GPI achieves a given contrast ratio sensitivity ∼ 50 times faster than the best previous-generation systems. A 600-star survey of young nearby stars will begin in 2014, with the goal of producing a sample of directly imaged planets that: (a) spans a broad range of temperatures, ages, and masses, opening up new areas of planetary atmosphere phase space; and (b) probes the range of semi-major-axes and stellar ages inaccessible to Doppler and transit surveys to produce statistical constraints on planet-formation model. Simulations predict that this survey will discover 25-50 exoplanets [16] with masses as low as that of Jupiter and at separations as close as 3-5 AU. Using a polarimetry mode, the survey will also map debris disks. GPI will be available to the astronomical community for guest observer projects, ranging from studying young solar system objects to the outflows from evolved stars. With the advent of GPI and similar dedicated facilities, high contrast imaging, spectroscopy and polarimetry will open up a new segment of other solar systems to characterization.   Table 3. Keplerian orbital elements. The quoted value is the median (50th percentile) of the marginalized posterior distributions with errors representing the 68% confidence interval (16 and 84th percentiles.) Derived quantitiesperiod, P , periapse and apoapse distances are also listed.   represented degeneracy in ω, e.g., see ω versus a, is good evidence of reliable sampling of the posterior distribution. Vertical dotted lines in the histograms denote the 68% confidence interval.