DIFFRACTION-LIMITED VISIBLE LIGHT IMAGES OF ORION TRAPEZIUM CLUSTER WITH THE MAGELLAN ADAPTIVE SECONDARY ADAPTIVE OPTICS SYSTEM (MagAO)*

L. M. Close; J. R. Males; K. Morzinski; D. Kopon; K. Follette; T. J. Rodigas; P. Hinz; Y-L. Wu; A. Puglisi; S. Esposito; A. Riccardi; E. Pinna; M. Xompero; R. Briguglio; A Uomoto; T. Hare

doi:10.1088/0004-637X/774/2/94

1. INTRODUCTION

1.1. The Need for High-resolution Imaging

It is critical to the understanding of the motions and masses of stars, brown dwarfs, and exoplanets to obtain the highest resolution images possible. In fact, almost every aspect of astronomical science benefits from the highest spatial resolutions possible. The highest resolution "maps" at the milliarcsec (mas) scales resolution (0 farcs 001 = 1 mas) are being produced by interferometry (like VLTI/AMBER in the IR). However, interferometric techniques suffer from incomplete uv coverage and object models are usually required to interpret interferometric data. Moreover, combining multiple 8 m telescopes together in the VLTI and waiting for the Earth's rotation to fill in the uv plane is both time-consuming and expensive (hence limiting the general utility of large surveys with VLTI, for example).

Imaging from space with a filled aperture (and so complete uv coverage) with the Hubble Space Telescope (HST) has proven to be very productive, but HST's small 2.4 m aperture, combined with a need for large pixels, limits its best spatial resolutions to 50–100 mas. Also, HST is considerably more expensive than any other telescope and its lifetime is limited. Large (8–10 m) ground-based telescopes can match HST's ∼ 50 mas V-band resolution with adaptive optics in the NIR (1.2–2.4 μm). For example, the 8.4 m LBT AO system (FLAO; Esposito et al. 2011) can also achieve deep 50 mas resolution images with AO at 1.64 μm (Close et al. 2012b). However, to achieve deep images better than ∼40–50 mas is impossible with the current generation of facility AO systems and cameras. For example, to reach ⩽20 mas resolutions at H band (1.65 μm) would require a D ⩾ 16 m filled-aperture telescope. Hence, it will not be until the ELT era (early to mid-2020s) that images in the NIR will be significantly sharper than 20 mas.

1.2. Into the Blue: Adaptive Optics in the Visible

However, there is another approach to reaching these resolutions. While 8–10 m AO performance is limited to ⩾40 mas in the NIR, it is possible to gain a factor of two improvement in spatial resolution by moving to shorter (bluer) wavelengths for AO correction. This so-called "visible AO" can theoretically reach 16 mas resolutions on an 8 m telescope at 0.656 μm (H_α). However, the complexity of an 8–10 m class AO system designed for optical wavelengths (>500 modes, at >1 kHz) is beyond that of the current facility systems (with perhaps the exception of the FLAO system on the 8.4 m LBT which, however, currently has no facility visible AO CCD science camera; Esposito et al. 2010).

We note that AO with "lucky" imaging in the visible has been successfully used at the somewhat smaller 5 m Palomar (Law et al. 2009) and has reached resolutions of 35 mas, and recently the Palm3000 system has demonstrated excellent corrections (Dekany et al. 2013). Improved Lucky visible imaging (image synthesis based on Fourier Amplitude selection) has also been developed by Garrel et al. (2012). Visible AO has been done before on much smaller telescopes like Robo-AO on the 1.5 m at Palomar (Baranec et al. 2012) or the Villages project on the 1.0m Nickel at Lick (Morzinski et al. 2010). In the near future, some polarization work will be done in the visible with the 8 m Very Large Telescope (VLT) with the SPHERE AO system and ZIMPOL (Bazzon et al. 2012). However, the Magellan Adaptive Optics system (MagAO) is the first large (D ⩾ 6.5 m) telescope AO system designed to work in the visible—complete with a facility CCD AO science camera (VisAO). The MagAO commissioning results presented here inform us on the utility of large telescope visible AO performance.

1.3. The Magellan AO System

We have developed an AO system (inspired, in large part, by LBT's FLAO system; Esposito et al. 2012) that can reach 20 mas resolutions with just 250–378 modes at 1 kHz sampling speeds. It is important that such a visible AO system be located at an excellent site where the median seeing is less than 0 farcs 64. To achieve an AO fitting error small enough to reach 110 nm rms total wavefront error (WFE) with 250 modes requires a telescope diameter of D ⩽ 6.5 m. Hence, a solution to this design problem is a fast (<1 ms response time) 585 element second generation adaptive secondary mirror (ASM) with a 1 kHz Pyramid wavefront sensor (PWFS). These are exactly the characteristics of the MagAO deployed on the 0 farcs 64 median seeing (Thomas-Osip et al. 2010) 6.5 m Magellan telescope at Las Campanas Observatory, Chile.

MagAO, with its VisAO camera,⁶ is the first large telescope (⩾6.5 m) facility AO system deployed that is targeting observations in the visible (0.6–1.1 μm). As will be shown in later in this paper, MagAO at first light produced long exposure (60 s) diffraction-limited (110 nm WFE; 28% Strehl) 0.63 μm images. Note that during the first light run (commissioning run 1; 2013 November/December) we were limited to 250 corrected modes.⁷ For more technical details about MagAO, please see Close et al. (2012a).

It is important to note that MagAO sends all the infrared light into the Clio2 NIR (1–5.3 μm) camera (Hinz et al. 2010; K. Morzinski et al. 2013, in preparation), whereas the visible light (λ < 1.1 μm) is split by a selectable beamsplitter between the PWFS and the VisAO (0.6–1.1 μm) science camera (for more on the VisAO camera, see Males et al. 2012; Kopon et al. 2013; Close et al. 2012a). Hence all three focal stations (Clio2, VisAO, and PWFS) simultaneously work on all targets, allowing visible and IR observations to be performed simultaneously.

1.4. First Light VisAO Science: Motions of the Massive Young Stars in the Orion Trapezium Cluster

Clearly, the exciting possibility of obtaining ∼20 mas FWHM images with MagAO could enhance our understanding of the positions (and motions) of the nearest massive young stars. Hence we targeted the Orion Trapezium cluster during the first light commissioning run with the MagAO system.

The study of the motions of the young stars in the Trapezium cluster is an important problem (see, for example, McCarthy & Zuckerman 2004; Close et al. 2012b; Grellmann et al. 2013). After all, the detailed formation of stars is still a poorly understood process. In particular, the formation mechanism of the lowest mass stars and brown dwarfs is uncertain. Detailed three-dimensional (and N-body) simulations of star formation byBate et al. (2002, 2003), Bate (2009, 2012), and Parker et al. (2011) all suggest that stellar embryos frequently form "mini clusters" which dynamically decay, "ejecting" the lowest mass members. Such theories can explain why there are far more field brown dwarfs compared to brown dwarf companions of solar-type stars (McCarthy & Zuckerman 2004) or early M stars (Hinz et al. 2002). Moreover, these theories which invoke some sort of dynamical decay (Durisen et al. 2001) or ejection (Reipurth & Clarke 2001) suggest that there should be no wide (>20 AU) very low mass (VLM; M_tot < 0.185 M_☉) binary systems observed in the field (age ∼5 Gyr). Indeed, the AO surveys of Close et al. (2003a) and the HST surveys of Reid et al. (2001), Burgasser et al. (2003), Bouy et al. (2003), and Gizis et al. (2000) have not discovered more than a few wide (>16 AU) VLM systems in the field population (for a review, see Burgasser et al. 2007). Additionally, the dynamical biasing toward the ejection of the lowest mass members naturally suggests that the frequency of field VLM binaries should be much lower (≲ 5% for M_tot ∼ 0.16 M_☉) than for more massive binaries (∼60% for M_tot ∼ 1 M_☉). Indeed, observations suggest that the binarity of VLM systems with M_tot ≲ 0.185 M_☉ is 10%–15% (Close et al. 2003a; Burgasser et al. 2003, 2007) which, although higher than predicted, is still lower than that of the ∼42% of more massive M-dwarfs (Fischer & Marcy 1992) or ∼60% of G star binaries (Duquennoy & Mayor 1991). However, as is noted in Close et al. (2007), there is evidence that in young clusters wide VLM binaries are much more common than in the old field population. They attribute this to observing these wide VLM systems before they are destroyed by encounters in their natal clusters. Hence, we need to look at nearby young clusters to see these low-mass objects in "mini clusters" (of a few bound stars) before ejection has occurred.

Despite the success of these decay, or ejection, scenarios in predicting the observed properties of low mass VLM stars and binaries, it is still not clear whether or not "mini clusters" even exist in the early stages of star formation. To better understand whether such "mini clusters" do exist, we have examined the closest major OB star formation cluster for signs of such "mini clusters." Here we focus on the θ¹ Ori stars in the famous Orion Trapezium cluster. Trying to determine if some of the tight star groups in the Trapezium cluster are gravitationally bound is a first step to determining if bound "mini clusters" exist. Also, it is important to understand the true number of real, physical, binaries in this cluster, as there is evidence that the overall number of binaries is lower (at least for the lower mass members) in the dense trapezium cluster compared to the lower density young associations like Taurus-Auriga (McCaughrean 2000; Kohler et al. 2006). In particular, we examine the case of the θ¹ Ori A, B and C groups in detail.

The Trapezium OB stars (θ¹ Ori A, B, C, D, and E; see Figure 1) consist of the most massive OB stars located at the center of the Orion Nebula star formation cluster (for a review, see Genzel & Stutzki 1989). Due to the nearby (Very Long Baseline Array trigonometric parallax distance of 414 ± 7 pc; Menten et al. 2007) and luminous nature of these stars, they are a unique laboratory in which to study a high-mass star formation cluster (the dominant birthplace for stars of all masses), and have been the target of several high-resolution imaging studies. For brevity, here we do not reproduce a complete history of past high resolution surveys of Trapezium; see Close et al. (2012b) instead for a review.

Close et al. (2012b) utilized the LBT FLAO system to map out the Trapezium in narrow-band NIR filters at ∼50–60 mas resolutions. In total, Close et al. (2012b) analyzed 14 yr of observations of the cluster. However, only the LBT 2011 observations were of very high quality. In this paper, we present the first high-resolution visible (0.57–0.68 μm) AO images. These images are of the Trapezium cluster and reach very high resolutions of ∼23 mas. We have now over 15 yr of observations of this field with at <0 farcs 08 resolution. More importantly, we now have two complete high-quality datasets from LBT and MagAO that track the motion of the Trapezium stars at <0 farcs 05 resolutions.

In this paper, we outline how these MagAO observations were carried out with the new VisAO camera. We detail how these data were calibrated and reduced and how the stellar positions were measured. We resolve the 32 mas binary θ¹ Ori C in a filled aperture image for the first time. We compare the measured astrometry for θ¹ Ori C1 and C2 against its published (interferometric) orbit. We also fit the observed positions to calculate velocities (or upper limits) for the θ¹ Ori B₁, B₂, B₃, B₄ and A₁, A₂ stars. While Schertl et al. (2003) and Close et al. (2003b, 2012b) hinted that the θ¹ Ori B group may be a bound "mini cluster," we show that here it is clearly so, with the first detection of curvature in the orbital motion of members of this group. We also present the first model for how the complex set of orbits in the θ¹ Ori B mini cluster could (and cannot) be arranged.

2. INSTRUMENTAL SET-UP

We utilized MagAO to obtain the first diffraction-limited (and unsaturated) images of the young stars in the Trapezium cluster in the visible (0.6–0.7 μm). This is not a simple task, since, as telescopes have increased in size, bright stars now tend to saturate even in the shortest possible exposures. Hence, special precautions are needed to avoid saturation of the bright Trapezium stars themselves. It is difficult to make unsaturated, but diffraction-limited, "visible light" images of the bright Trapezium stars with modern 6.5 m class AO systems at even moderately high Strehl. Note that this is the first such dataset ever published. The following subsections outline how this was accomplished.

The MagAO system is unique (at least in the southern hemisphere) in many ways. To reduce the aberrations caused by atmospheric turbulence, all large telescope AO systems have a deformable mirror (DM) which is updated in shape at ∼500 Hz. Except for the MMT AO and LBT AO systems (Wildi et al. 2003; Esposito et al. 2011), all other adaptive optics systems have located this DM at a reimaged pupil (effectively a compressed image of the primary mirror). To reimage the pupil onto a DM typically requires an additional three to eight warm optical surfaces, which significantly increases the thermal background and decreases the optical throughput of the system (Lloyd-Hart 2000). However, MagAO utilizes a next generation adaptive secondary DM. This DM is both the secondary mirror of the telescope and the DM of the AO system. In this manner, there are no additional optics required in front of the science camera. Hence, the emissivity is lower, and the throughput is higher. MagAO's DM is an advanced "second generation" ASM (similar to those on the LBT), which enables the highest on-sky visible Strehl (>25% at r' band; 0.57–0.68 μm) of any large 6.5–10 m telescope today.

The MagAO ASM consists of 585 voice coil actuators that push (or pull) on 585 small magnets glued to the backsurface of a thin (1.6 mm), 850 mm aspheric ellipsoidal Zerodur glass "shell" (for a detailed review of the secondary mirror, see Close et al. 2012a). As in the case of the LBT AO system, we have complete positional control of the surface of this reflective shell by use of a 70 kHz capacitive sensor feedback loop. This positional feedback loop allows one to position an actuator of the deformable shell to within ∼5 nm rms (total residual polishing WFEs (mainly at interactuator scales) amount to only ∼50 nm rms over the whole secondary). The AO system samples (and drives the ASM) at 990 Hz using 250-378 active controlled modes (with 585 actuators) on bright stars (R < 9 mag).⁸

The wavefront slopes are measured with a very accurate, well calibrated, low aliasing error PWFS. This is the second large telescope to use a PWFS (after the LBT; Esposito et al. 2011). The performance of the MagAO PWFS is excellent. The very low residual WFEs obtained by the PWFS + ASM combination is due in part to the very accurate (high signal to noise) interaction matrix that can be obtained in closed-loop daytime calibrations with a retro-reflecting calibration return optic (CRO; Kopon et al. 2013) that takes advantage of the Gregorian (concave) nature of the secondary. To guarantee strict "on-sky" compliance with the "daytime calibrated" interaction matrix pupil/ASM/PWFS geometry the PWFS utilizes a novel "closed-loop pupil alignment system" that maintains the pupil alignment to <2.5 μm (at the PWFS CCD39 images of the four pupils produced by the PWFS) during all closed-loop operations on bright stars. Moreover, we use a fixed pupil mask on the ASM to maintain the exact same pupil illumination when the CRO is used and also when we are on sky—so that our interaction matrices are valid (on and off sky). For a detailed review of the MagAO system, see Close et al. (2012a) and references within.

2.1. The MagAO PSF and Calculating Strehl

During the MagAO first light commissioning run, we observed the θ¹ Ori A, B and C groups on the nights of 2012 December 3, 4, and 8 (UT). The AO system corrected the lowest 250 system modes and was updated at 990 Hz. The PWFS pupil was close-loop stabilized and the shell was protected from wind with a windscreen at the secondary mirror. Cooling pumps (for Clio2, VisAO, and the PWFS) added some vibrational blurring into the point-spread function (PSF). After commissioning run 1 these pumps were much better isolated. Nevertheless, the PSFs were still close to perfectly diffraction-limited. To better gauge the effectiveness of the AO correction, we need to be able to measure the long-exposure PSF and calculate the Strehl of the PSF.

On bright (R < 9) guide stars in ∼0 farcs 6 V-band seeing we could obtain deep five-minute PSF images (with no SAA or post-detection processing) with Strehls of 43% at Y_short (Ys; 0.98 μm), or 140 nm rms WFE (by use of the extended Marechal's approximation; see Figure 2). We note the deep five-minute image in Figure 2 suffered from some additional vibrational blurring due to the cooling pumps for the CCDs and Clio2.⁹ These deep PSF images helped model the PSF to calculate Strehls for MagAO on θ¹ Ori C which was so bright that only a 64 × 64 CCD window could be readout without saturation on C₁. Hence the wings of the PSF (beyond the 64 × 64 window) had to be estimated from a wavelength scaled PSF "halo" model based on the measured deep PSF wings of Figure 2. In this manner, realistic Strehls could be estimated reliably from the small 64 × 64 images of θ¹ Ori C₁.¹⁰ We note that it was only the Strehl of θ¹ Ori C₁ that required this bootstrap approach; all other Strehls (from full frame CCD images) in this paper were measured in the usual manner by comparison to our model theoretical PSF.

**Figure 2.** Radial profile (red points) of a deep (300 s) MagAO Y_short (0.98 μm) PSF on a bright (V = 5 mag) star closed-loop at 990 Hz with 250 modes in 06 V band seeing (from J. R. Males et al. 2013, in preparation). Inset: a log10 Stretch of the PSF. There was no post-detection processing of any of the data (no SAA, no Lucky imaging, or frame selection applied). The theoretical MagAO PSF profile as imaged by the E2V CCD47 (Strehl 100%). A detailed comparison of the observed PSF to theory with our CCD47 (including dark current and PRF) shows that we reached a Strehl of 43% or 140 nm rms optical wavefront error.
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2.2. The VisAO CCD AO Science Camera

These observations utilized the first facility visible light AO science camera (VisAO; Males et al. 2012; Kopon et al. 2013). VisAO has a fast, frame transfer, 1024 × 1024 0.5–1.1 μm E2V CCD47 detector. We used the 64 × 64 window mode to minimize saturation on the array while observing θ¹ Ori C1 (V = 5.13 mag), while the ∼10 × fainter θ¹ Ori B1 (V = 7.2) allowed the whole CCD to be readout without saturation.

The VisAO focal plane platescales were calibrated by the astrometry of four stars in the HD 40887 quadruple system and θ¹ Ori B₁ and θ¹ Ori E₁¹¹ (see Sections 3 and 4 for more details about how the images were first reduced).

The positions (found by the IRAF allstar PSF fitting task) of these stars from our VisAO images were compared to unsaturated astrometry from Close et al. (2012b) which itself is derived from the HST Advanced Camera for Surveys (ACS) astrometry of Ricci et al. (2007). VisAO platescales and rms errors were then determined for the H_α, [O i], r', i', z' and Ys filters with the IRAF geomap task. The platescale found was 0 farcs 0078513 ± 0 farcs 000015 pix⁻¹ at H_α (0.656 μm), and [O i] (0.63 μm) providing a 8 farcs 03 × 8 farcs 03 field of view (FOV) with our f/52.5 beam on the CCD47's 13.0 μm pixels. At r' (0.63 μm) the platescale was slightly coarser at 0 farcs 007917 ± 0 farcs 000015 pix⁻¹, at z' (0.906 μm) just slightly finer at 0 farcs 007911 ± 0 farcs 000012 pix⁻¹, and at Y_short (0.982 μm) 0 farcs 007906 ± 0 farcs 000014 pix⁻¹. By design, the f/16 beam (direct from the ASM) is slowed down to f/52.5 to yield these very fine 7.9 mas pixel⁻¹ VisAO platescales. We note that this is one of the finest platescales ever for a facility camera.

Small distortions were detected by dithering a binary across the VisAO CCD. In this manner, we detected a small change in the Y platescale (⩽1%) from the top of the array to the bottom. The exact formula to correct a binary with a primary star at position X,Y of separation δx and δy for any residual distortions is true_δx = measured_δx − δdx/(abs[measured_δx]/110.0) and true_δy = measured_δy − δdy/(abs[measured_δy]/44.5) where δdx = −0.00038921676*(X − 512) + 0.00084322443*(Y − 512) and δdy = −0.00025760395*(X − 512) − 0.0024045175*(Y − 512). Our observations were near the center of the detector and so these corrections were actually very small (0.1%–0.5% or 0.1–5 mas changes to the 0 farcs 1–1'' binaries); nevertheless all binary observations in this paper have been fully distortion corrected.

To determine the orientation of the Y axis of the VisAO images (which were all taken with the rotator following) it was first necessary to rotate each image counterclockwise (with the IRAF rotate task) by the ROTOFF FITS keyword value +90 deg. At this point it was found by geomap that the direction of north was slightly (0 fdg 890) east of VisAO's Y axis compared to the HST ACS Ricci et al. (2007) and LBT images (Close et al. 2012b) of the field. Hence a final counterclockwise rotation of −0 fdg 890 was applied to the final image. The rms uncertainty adopted for the MagAO rotator angle is estimated as ∼0 fdg 3 this is the maximum error seen between different images of these stars on different nights. We suspect that this value of ∼0 fdg 3 is quite conservative based on the very low scatter in the PA fits shown later in this paper.

3. OBSERVATIONS AND REDUCTIONS

For the θ¹ Ori C field we locked the AO system (at 990 Hz, 250 modes) in 0 farcs 5–0 farcs 7 seeing on the bright O5pv binary star θ¹ Ori C (V = 5.13 mag) and used a 64 × 64 window in the center of the VisAO CCD with a set of 2608 × 0.023 s (60 s total) unsaturated exposures at H_α, [O i], and r'. Immediately following the unsaturated exposures, a set of 60 s exposures were obtained with the AO off. We note that θ¹ Ori C is really a ∼0 farcs 03 binary composed of C1 and C2 (see Kraus et al. 2007 for more details).

Then the AO system was locked (250 modes, 990 Hz) on θ¹ Ori B1 (V = 7.96 mag) and VisAO was used over its fullframe (1k × 1k) pixels to produce a set of 212 × 0.283 s unsaturated (60 s total) images at z'.

The individual frames were reduced in a normal manner. We used our custom AO image reduction script of Close et al. (2003a) to sky/bias subtract, cross-correlate (when needed), and median combine each image. The final individual image sets of the C and B fields each had a total exposure time of 1 minute. Figure 1 is a large FOV LBT NIR AO image from Close et al. (2012b) that defines the nomenclature and relative positions of the Trapezium stars for clarity.

In Figure 2, one can see the marked improvement in resolution (∼600 mas–34 mas) and Strehl (∼0.5%–43%) having the AO loop closed makes to a 300 s exposure. We note these Ys images are not post-detection "frame selected" (lucky imaging) nor shift and added (SAA), so they are true 300 s open-shutter exposures.

In Figure 3, we show typical images of the binary θ¹ Ori C₁ and C₂ imaged in 0 farcs 5 seeing ([O i] and r') on 2012 December 8 and worse 0 farcs 7 seeing for H_α on 2012 December 3. In all cases, excellent (26–29 mas and 28%–25% Strehl) images are obtained.

In the middle row of Figure 3, we retrieve the true resolution of the optical beam on the CCD47 by post-detection alignment of the images (∼2–4 mas improvement). We also make a similarly small resolution improvement (∼2–4 mas) by removing the blurring effects of the CCD47's pixel response function (PRF). The CCD47's 13.0 μm pixel PRF was calibrated by noting the slight improvement in FWHM when a lab CCD with smaller 5.5 μm pixels were used (instead of the 13 μm pixels) in a PRF lab test. A similar amount (just slightly less) PRF is observed with HST's ACS CCDs.¹² Once vibrations and PRF are minimized, the images have 21–23 mas resolutions.

Very short (23 ms) individual images were not effected as much by the residual vibrations and achieved very high resolutions of 21 mas (see Figure 4). These vibrations were found in commissioning to be mainly residual 60 Hz vibrations not corrected by MagAO and are likely due to a few fans on the telescope that we could not turn off. However, once corrected for PRF, these images are diffraction-limited (FWHM = 19 mas; Strehl = 54%; see Figure 4). We do not use Lucky imaging in this study, since the long exposure (60 s) images in Figure 3 are much deeper (and almost as sharp) as those possible to obtain with Lucky in 60 s of telescope time.

4. ASTROMETRY AND PHOTOMETRY

All reduced (with SAA but not PRF-corrected) 60 s images of θ¹ Ori C₁C₂ were analyzed with the DAOPHOT PSF fitting task allstar (Stetson 1987). The ±0.48 mas astrometric error of this very tight binary (where our ∼0.03 mas platescale errors can be ignored) was estimated by the standard deviation of the astrometry differences between the three filters ([O i], r', Hα¹³) used. Our θ¹ Ori C₁C₂ measurements of 32.64 ± 0.48 and PA = 206 fdg 31 ± 0 fdg 17 are compared to the interferometrically derived orbit in Figure 5. We find reasonable agreement between the AO images and the interferometrically derived orbit of Kraus et al. (2009). For θ¹ Ori B₁, B₂, B₃, B₄ and θ¹ Ori A₁, A₂ the astrometry are summarized in Table 1.

**Figure 5.** PSF fitting photometry of the [O i], r', and H_α images in Figure 3 of θ¹ Ori C₁C₂ with a 44% Strehl theoretical PSF gave a good fit to both binary components in all cases. For 2012 December 12 UT we find separation is 32.64 ± 0.48 mas and PA is 20631 ± 017. Here we plot this position (in red) against the interferometric orbit of Kraus et al. (2009). The agreement with the predicted position of C₂ w.r.t. C₁ is reasonable given the uncertainty of the orbital solution. However, more VisAO astrometry at this level of accuracy following this poorly sampled side (20° < PA < 210°) of the orbit over the next few years would certainly lead to a better orbital solution than is known today.
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fdg — **Figure 5.** PSF fitting photometry of the [O i], r', and H_α images in Figure 3 of θ¹ Ori C₁C₂ with a 44% Strehl theoretical PSF gave a good fit to both binary components in all cases. For 2012 December 12 UT we find separation is 32.64 ± 0.48 mas and PA is 20631 ± 017. Here we plot this position (in red) against the interferometric orbit of Kraus et al. (2009). The agreement with the predicted position of C₂ w.r.t. C₁ is reasonable given the uncertainty of the orbital solution. However, more VisAO astrometry at this level of accuracy following this poorly sampled side (20° < PA < 210°) of the orbit over the next few years would certainly lead to a better orbital solution than is known today.
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Table 1. High Resolution Observations of the θ¹ Ori B and A Groups

System	ΔH	ΔK'	Separation	Sep. Vel.	PA	PA Vel.	Telescope	Epoch
Name	(mag)	(mag)	('')	(Sep. mas yr⁻¹)	(°)	(° yr⁻¹)	Telescope	(m/d/y)
B₁B₂	2.30 ± 0.15		0.942 ± 0.020		254.9 ± 1.0		SAO^a	10/14/97
		1.31 ± 0.10^b	0.942 ± 0.020		254.4 ± 1.0		SAO^a	11/03/98
		2.07 ± 0.05	0.9388 ± 0.0040		254.6 ± 1.0		GEMINI	09/19/01
	2.24 ± 0.05		0.9375 ± 0.0030		255.1 ± 1.0		MMT	01/20/03
			0.9411 ± 0.0023		254.55 ± 0.3		LBT	10/16/11
			0.9415 ± 0.0014		254.64 ± 0.3		MagAO	12/04/12
			with MagAO=	0.31 ± 0.25		−0.009 ± 0.043
			corr.=	89%; no vel.	corr.=	33%; no vel.
				detected^c		detected^c

B₂B₃	1.00 ± 0.11		0.114 ± 0.05		204.3 ± 4.0		SAO^a	10/14/97
		1.24 ± 0.20	0.117 ± 0.005		205.7 ± 4.0		SAO^a	11/03/98
		1.04 ± 0.05	0.1166 ± 0.0040		207.8 ± 1.0		GEMINI	09/19/01
	0.85 ± 0.05		0.1182 ± 0.0030		209.7 ± 1.0		MMT	01/20/03
			0.1156 ± 0.0005		220.39 ± 0.3		LBT	10/16/11
			0.1160 ± 0.0002		221.50 ± 0.3		MagAO	12/04/12
			with MagAO=	−0.04 ± 0.14		1.19 ± 0.06
			corr.=	24%; no vel.		4.7 ± 0.2 km s⁻¹
				detected		corr.=99.9%

B₁B₄		5.05 ± 0.8	0.609 ± 0.008		298.0 ± 2.0		SAO^d	02/07/01
		5.01 ± 0.10	0.6126 ± 0.0040		298.2 ± 1.0		GEMINI	09/19/01
	4.98 ± 0.10		0.6090 ± 0.0050		298.4 ± 1.0		MMT	01/20/03
			0.6157 ± 0.003		300.1 ± 0.5		LBT	10/16/11
			0.6182 ± 0.0009		300.23 ± 0.3		MagAO	12/04/12
			with MagAO=	0.72 ± 0.23		0.181 ± 0.067
				1.4 ± 0.5 km s⁻¹		3.83 ± 1.27 km s⁻¹
				corr.=95%		corr.=99.9%

A₁A₂	1.51 ± 0.15	1.38 ± 0.10	0.208 ± 0.030		343.5 ± 5.0		Calar Alto^e	11/15/94
		1.51 ± 0.05	0.2215 ± 0.005		353.8 ± 2.0		SAO^a	11/03/98
		1.62 ± 0.05	0.2051 ± 0.0030		356.9 ± 1.0		GEMINI	09/19/01
			0.1931 ± 0.0005		366.5 ± 0.3		LBT	10/16/11
			0.1881 ± 0.0016		367.6 ± 0.3		MagAO^f	12/08/12
			with MagAO=	−1.6 ± 0.2		0.98 ± 0.07
				−3.2 ± 0.3 km s⁻¹		6.3 ± 0.4 km s⁻¹
				corr.= 92.9%		corr.= 99.4%

Notes. ^aSpeckle observations of Weigelt et al. (1999). ^bThese low ΔK values are possibly due to θ¹ Ori B₁ being in eclipse during the 11/03/98 observations of Weigelt et al. (1999). ^cNote there is velocity detected from B₁ w.r.t. the barycenter of the B₂B₃ binary; see Figures 10 and 11. ^dSpeckle observations of Schertl et al. (2003). ^eSpeckle observations of Petr et al. (1998). ^fA₁A₂ Data from Ks image from the MagAO/Clio2 NIR camera (K. Morzinski et al. 2013, in preparation).

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In the θ¹ Ori B group the PSF star used was the unsaturated θ¹ Ori B₁ itself. Since all the members of the θ¹ Ori B group are located within 1'' of θ¹ Ori B₁, the PSF fit is particularly excellent there (there is no detectable change in PSF morphology due to anisoplanatic effects inside the θ¹ Ori B group; Diolaiti et al. 2000). Moreover, the residuals over the whole field were less than a few percent after PSF subtraction. This is not really surprising given the quality of the nights combined with the fact that no star was farther than ∼1'' from the guide star. However, to minimize this affect, we only used the longer wavelength z' images reduced with SAA (taken on 2012 December 4) where anisoplanatic PSF effects were undetected. The relative positional accuracy is an excellent ∼0.2–1.4 mas in radial separation. The ∼0.2–1.4 mas separation errors are the resultant of the platescale uncertainty added in quadrature with the measurement uncertainty (FWHM/(S/N)). The errors are somewhat worse in the PA direction (0.6–5 mas) due to a fixed ±0.3 deg conservative estimate of our absolute rotator uncertainly.

We can also compare our MagAO data to older (somewhat less accurate) images of the Trapezium B stars from Close et al. (2003b) who used AO images from Gemini and the 6.5 m MMT and speckle images from the literature (Schertl et al. 2003). Even though these individual observations are of lower quality and Strehl than the MagAO ones (compare Figures 6–8 to that of MagAO in Figure 9), the 15 years between these observations and those of MagAO can highlight even very small orbital motions of bound systems in the Trapezium. It also shows the very significant improvement in high Strehl AO now possible with PWFS and next generation ASMs.

**Figure 6.** 8 m Gemini/Hokupa'a images of the θ¹ Ori B group in the K' band (2001 September 19; from Close et al. 2003b). Resolution 0085. Log scale. North is up and east is to the left.
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**Figure 7.** Detail of the θ¹ Ori B group as imaged at 0077 (Strehl >20%) resolution (in the H band) with the MMT AO system (2003 January 20) from Close et al. (2003b).
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**Figure 8.** LBT AO Brγ (2.16 μm) images of the θ¹ Ori B group. Resolution 006. Logarithmic color scale. North is up and east is to the left. Strehl is ∼75% (from Close et al. 2012b).
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**Figure 9.** MagAO z' (0.91 μm) images of the θ¹ Ori B group. Resolution 0034. Linear color scale. North is up and east is to the left. Note that this image is ∼2 × sharper than that of the 8.4 m LBT at Brγ (2.16 μm; see Figure 8). This is clear evidence that AO in the "blue" allows a smaller 6.5 m telescope to outperform the resolution of an 8 m in the NIR.
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A test to see how accurate our astrometry is over the last 15 yr is to look at the scatter from a linear trend of the θ¹ Ori B group's motions. A comparison of our highly accurate positions with the historical positions from the literature is summarized in Table 1. Linear (weighted by astrometric error) fits to the data in Table 1 (Figures 10–15) yield the velocities shown in Table 1. The overall error in the relative proper motions is now an impressive ≲ 0.2 mas yr⁻¹ in proper motion (≲ 0.4 km s⁻¹), a factor of 2 improvement in accuracy when the VisAO positions are added into these calculations, compared to the last published values from Close et al. (2012b).

**Figure 10.** Separation between θ¹ Ori B₁ and the barycenter of the B₂B₃ binary. Note how over 15 yr of observation there has been a small, yet significant, relative proper motion observed (0.80 ± 0.18 mas yr⁻¹; which is a significant correlation at the 97.4% level). The first two data points are speckle observations from the 6 m SAO telescope (Weigelt et al. 1999), the next point is from Gemini/Hokupa'a observations Close et al. (2003b) followed by MMT AO observations Close et al. (2003b), and then LBT Close et al. (2012b), and the last from the MagAO system (this study).
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**Figure 11.** Position angle between θ¹ Ori B₁ and the barycenter of the B₂B₃ binary. Note how over 15 yr of observation there has been little relative PA motion observed (0030 ± 0044 yr⁻¹ which is just significant at the 88% level). The epochs of the data are the same as in Figure 10.
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fdg — **Figure 11.** Position angle between θ¹ Ori B₁ and the barycenter of the B₂B₃ binary. Note how over 15 yr of observation there has been little relative PA motion observed (0030 ± 0044 yr⁻¹ which is just significant at the 88% level). The epochs of the data are the same as in Figure 10.
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**Figure 12.** Separation between the θ¹ Ori B₂ and B₃ components. Note the surprising lack of any significant relative motion (−0.04 ± 0.14 mas yr⁻¹). The rms scatter from a constant value is only 0.14 mas yr⁻¹. There appears to be very little change in the separation of the B₂B₃ system. The epochs of the data are the same as in Figure 10.
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**Figure 13.** Position angle of θ¹ Ori B₂ and B₃. Here we clearly observe the real orbital arc of motion where B₃ is moving counterclockwise (at 1.19 ± 006 yr⁻¹; a correlation significant at the 99.9% level) around B₂. The epochs of the data are the same as in Figure 10.
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fdg — **Figure 13.** Position angle of θ¹ Ori B₂ and B₃. Here we clearly observe the real orbital arc of motion where B₃ is moving counterclockwise (at 1.19 ± 006 yr⁻¹; a correlation significant at the 99.9% level) around B₂. The epochs of the data are the same as in Figure 10.
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**Figure 14.** Separation between θ¹ Ori B₁ and B₄. Note how over 13 yr of observation there is little change in the separation (0.72 ± 0.23 mas yr⁻¹; correlation 95%). The first data point is an speckle observation from the 6 m SAO telescope (Schertl et al. 2003), then Gemini/Hokupa'a observation (Close et al. 2003b), the next data point is from the MMT AO observation (Close et al. 2003b), the next from the LBT (Close et al. 2012b), and the last is from MagAO.
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**Figure 15.** Position angle between θ¹ Ori B₁ and B₄. Note how over 13 yr of observation there has been only now clear significant relative proper motion observed (0181 ± 0067 yr⁻¹; correlation 99.8%). The sources of the data are the same as in Figure 14.
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fdg — **Figure 15.** Position angle between θ¹ Ori B₁ and B₄. Note how over 13 yr of observation there has been only now clear significant relative proper motion observed (0181 ± 0067 yr⁻¹; correlation 99.8%). The sources of the data are the same as in Figure 14.
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5. ANALYSIS AND DISCUSSION

With these accuracies it is now possible to determine whether these stars in the θ¹ Ori B group are bound together, or merely chance projections in this very crowded region. We adopt the masses of each star from the Siess Forestini & Dougados (1997); Bernasconi & Maeder (1996) tracks fit by Weigelt et al. (1999) where we find masses of: B₁ ∼ 7 M_☉, B₂ ∼ 3 M_☉, B₃ ∼ 2.5 M_☉, B₄ ∼ 0.2 M_☉, B₅ ∼ 7 M_☉, A₁ ∼ 20 M_☉, A₂ ∼ 4 M_☉, and A₃ ∼ 2.6 M_☉. Based on these masses (which are similar to those adopted by Schertl et al. 2003) we can comment on whether the observed motions are less than the escape velocities expected for simple face-on circular orbits.

Our combination of high spatial resolution and high signal to noise shows that there is very little significant motion in the B₁B₂ system over the last 15 years (as we might expect since the rotation of B₂ about the barycenter of the B₂B₃ system appears to be just canceling the motion of B₂ w.r.t. B₁). But of course, it is really the barycenter of the tight B₂B₃ binary that would be in orbit around B₁. Hence the barycenter would have to show steady orbital motion if bound to B₁. Since B₂ is only 20% more massive than B₃, this means the B₂B₃ barycenter is currently 52 mas at PA 221 fdg 5 from the center of B₂. In Figures 10 and 11, we see that there is a small yet significant motion of the barycenter of B₂B₃ w.r.t. B₁ of some 0.80 ± 0.18 mas yr⁻¹ (1.6 ± 0.3 km s⁻¹) and in PA by 0 fdg 030 ± 0 fdg 044 yr⁻¹ (1.0 ± 1.0 km s⁻¹). Hence the motion of B₂B₃ is currently about 1.9 ± 0.6 km s⁻¹ in the direction of PA ∼305° (moving toward the WNW direction from B₁). This is the first time this motion has been detected. At this time it is not yet possible to prove that this is true orbital motion, but, given how close B₂B₃ is to B₁, this is likely orbital motion.

We have, of course, observed clear orbital motion (at 4.7 ± 0.2 km s⁻¹) in the very tight θ¹ Ori B₂B₃ system in almost pure PA (see Figure 13). In fact, now that we have observed over 20° of PA rotation (with no significant change in separation), we have clear evidence of an "arc" of the curvature of the system. The motion of the B₂B₃ binary is roughly consistent with a face-on, circular orbit (orbiting in the counterclockwise direction). A mildly elliptical orbit is also quite plausible given the very small amount of orbital phase observed to date.

We also see linear orbital motion of 7.0 ± 0.5 km s⁻¹ in the θ¹ Ori A₁A₂ system (see Table 1). This is consistent with the motion seen by Grellmann et al. (2013). We know that this is likely orbital motion since it is higher than the motion of unrelated stars in the cluster due to their very close separation of just 0 farcs 19.

5.1. Is the θ¹ Ori B₂B₃ System Physical?

The relative velocity in the θ¹ Ori B₂B₃ system (in the plane of the sky) is now more accurate by ∼10 × compared to that of Close et al. (2003b) and by ∼2 × compared to Close et al. (2012b). Our new velocity of 4.7 ± 0.2 km s⁻¹ is consistent, but with much lower errors, with the ∼4.2 ± 2.1 km s⁻¹ of (Close et al. 2003b; this velocity is in the azimuthal direction; see Figure 13). This is a reasonable V_tan since an orbital velocity of ∼6.7 km s⁻¹ is expected from a face-on circular orbit from a ∼5.5 M_☉ binary system like θ¹ Ori B₂B₃ with a 49 AU projected separation (implying an orbital period of the order of ∼200 yr). This theoretical value of ∼200 yr is close to the 302 ± 16 yr orbit that comes from assuming the current measured angular (PA) velocity stays constant. It is worth noting that this velocity is also greater than the ∼3 km s⁻¹ Hillenbrand & Hartmann (1998) velocity dispersion of the cluster.

Our observed velocity of 1 fdg 19 ± 0 fdg 06 yr⁻¹ is ∼10 × more accurate than that of Close et al. (2003b). This primarily azimuthal motion strongly suggests a curving orbital arc of B₃ orbiting B₂ counterclockwise.

5.2. Is the θ¹ Ori B Group Stable Long Term?

The barycenter of B₂B₃ is moving at a low 1.9 ± 0.6 km s⁻¹ in the plane of the sky w.r.t. to B₁ (itself a very tight pair with B₅) where the escape velocity V_esc ∼ 6 km s⁻¹ for this massive system (∼20 M_☉). Hence these two pairs are likely gravitationally bound together. This is the first effort that has measured this small barycenter velocity definitively. Hence, we can say that these two pairs currently form a rare bound "mini cluster" of young massive stars.

5.2.1. Is the Orbit of θ¹ Ori B₄ Stable?

The two AO measurements of (Close et al. 2003b; and the one speckle detection of Schertl et al. 2003) did not detect a significant velocity of B₄ w.r.t. B₁: 2 ± 11 km s⁻¹. However, our much better data and timeline between the LBT epoch and the excellent MagAO observations has shed some light on the question of B4 orbiting B1. As is clear from Figures 14 and 15, there appears to be a real velocity of 4.1 ± 1.2 km s⁻¹ detected. This is greater than the random velocity of the cluster yet below the escape velocity of ∼6 km s⁻¹, and this points toward B₄ also being a gravitationally bound member of the θ¹ Ori B group. Again we are observing almost pure motion in PA (0 fdg 181 yr⁻¹ counterclockwise). Assuming a simple face-on orbit, we would expect a very rough period of ∼2000 ± 700 yr for B₄ to orbit B₁ given an average angular velocity of 0 fdg 181 yr⁻¹.

5.3. A Possible Model for the Orbits of the θ¹ Ori B Group

It is tempting to define a circular orbit baseline model of the θ¹ Ori B system with the center as the very tightly bound 6.47 day massive (∼14 M_☉) 0.13 AU spectroscopic binary B₁B₅ which we cannot spatially resolve. Around this center is the low mass 0.2 M_☉ B₄ some ∼254 AU away which orbits every ∼2000 ± 700 yr in a roughly face-on circular orbit. Then, farther out, the tight 49 AU binary B₂B₃ rotates every 302 ± 16 yr around its barycenter with roughly a face-on circular orbit. However, the co-planar geometry is broken by the orbit of this barycenter around B₁. It appears that the B₂B₃ barycenter is moving in a bound orbit to WNW (PA ∼305°). This motion cannot be in a simple face-on circular orbit, and so must be (if close to circular) inclined by about ∼30°, but many other elliptical orbits are also possible. We simply do not have enough of time baseline to understand the fine details of this orbit today. If we simply assume that it is a inclined circular orbit then it has roughly a ∼820 AU (deprojected) separation from B₁ and a period of some ∼11, 000 yr.

Refer to Figures 16 and 17 for illustrations of what these obits would look like if they are all close to circular. It is interesting to note that once the true (deprojected) separation of B₂B₃ is considered, the group seems more hierarchical than reported in Close et al. (2012b). For example, the ratio of the three main periods are P₂₃: P_1/4: P_1/23 = 1:7:36 so that there is almost an order of magnitude separating each period. This large spread of orbital periods will lend some stability to this "mini cluster." On the other hand, B₄'s VLM, its intermediate period, and its location w.r.t. to the other four groups members makes it highly unlikely that B₄ is on a long-term stable orbit within the group. It is very likely that an interaction between the much more massive B₂B₃ and B₄ will eject B₄ in the future—leading to a slightly more tightly bound "mini cluster" without B₄. As we will discuss in the next section, even the much more massive B₃ may not even be stable in the long-term.

**Figure 16.** Possible model of the motions of the θ¹ Ori B group. Here we show a *lucy* deconvolved image of the z' image. We make an assumption that all the orbits are circular and plot possible orbital solutions for each component's orbit about B₁ based on this rough assumption. We also plot the actual observed orbital "arcs" imaged so far over the last 15 yr for the system. Clearly the orbits are still undefined, but this plot gives some insight into the nature of the system.
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**Figure 17.** Zoom out of Figure 16 to show the full inclined orbit of the B₂B₃ barycenter around B₁.
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5.3.1. Is the Orbit of B₃ around B₂ and of B₅ around B₁ Stable in the Long-term?

Close et al. (2012b) noted that the distance $D_{B_1B_5} \sim 3 \times 10^{-4} \times D_{B_1B_5 B_2B_3}$ and thus the very tight (0.13 AU) tight B₁B₅ system is, of course, very stable. More interesting is the case of B₂B₃. Their deprojected distance is not very small compared to their projected distance (D) from the B₁B₅ pair: $D_{B_2B_3} \sim 0.06 \times D_{B_1B_5 B_2B_3}$ . Thus the stability of the B₂B₃ orbit needs a more detailed analysis since it is possible that B₃ may be ejected in the future.

The orbital period of the two binaries w.r.t. each other is P_1/23 ∼ 11, 000 yr, while the orbital period of B₃ w.r.t. B₂ amounts to P_2/3 ∼ 300 yr. For the calculation of both periods, we have assumed the masses as given above, and circular orbits in the plane of the sky (except for B₁B₃ which is inclined at ∼30°). This leads to a period ratio X = P_1/23/P_2/3 ∼ 36. Eggelton & Kiseleva's stability criterion requires X ⩾ X_crit = 10.08 for the masses in the B group. This means that within the accuracy limits of our investigation, the binary B₂B₃ is likely stable (different from the marginal stability found in Close et al. 2012b). The stability criterion depends also on the orbits' eccentricities. However, mild eccentricities of the order of e ∼ 0.1 (as can be expected to develop in hierarchical triple systems; see, e.g., Georgakarakos 2002) can make the B group unstable. However, the θ¹ Ori B system seems to be a good example of a highly dynamic star formation "mini cluster" which might, in the future, eject the lowest-mass member(s) through dynamical decay (Durisen et al. 2001).

6. CONCLUSIONS

In this study, we utilized the new high-order (585 actuator) MagAO to obtain very high-resolution science in the visible with MagAO's VisAO CCD camera. In the median seeing conditions of Magellan (0 farcs 5–0 farcs 7), we found that MagAO delivers individual short exposure images as good as 19 mas optical resolution. Due to residual 60 Hz vibrations, long exposure (60 s) r' (0.63 μm) images are slightly coarser at FWHM = 23–29 mas (Strehl ∼28%) with bright (R < 9 mag) guide stars. These are the highest resolution filled-aperture images published to date. Images of the young (∼1 Myr) Orion Trapezium θ¹ Ori A, B, and C cluster members were obtained with the VisAO camera. In particular, the 32 mas binary θ¹ Ori C₁C₂ was easily resolved in non-interferometric images for the first time. Relative positions of the bright trapezium binary stars were measured with 0.6–5.0 mas accuracy. We now are sensitive to relative proper motions of just ∼0.2 mas yr⁻¹ (∼0.4 km s⁻¹ at 414 pc)—this is a ∼2–10 × improvement in velocity accuracy compared to previous efforts. We now detect clear orbital motions of θ¹ Ori B₂B₃ and A₁A₂ of 4.7 ± 0.2 km s⁻¹ and 7.1 ± 0.5 km s⁻¹, respectively. For the first time, we see clear motion of the barycenter of θ¹ Ori B₂B₃ in about θ¹ Ori B₁. All five members of the θ¹ Ori B system appear likely a gravitationally bound "mini cluster," but we find that not all the orbits can be both circular and co-planar. The very lowest mass member of the θ¹ Ori B system (B₄; mass ∼0.2 M_☉) has a very clearly detected motion (at 4.1 ± 1.3 km s⁻¹; correlation = 99.9%) w.r.t. B₁. Previous work has suggested that B₄ and B₃ are both on long-term unstable orbits and will be ejected from this "mini cluster." However, our new "baseline" model of the θ¹ Ori B system suggests a more hierarchical system than previously thought, and so the ejection of B₄ may not occur for many orbits, and B₃ may be stable against ejection long-term. This "ejection" process of the lowest mass member of a "mini cluster" could play a major role in the formation of low mass stars and brown dwarfs.

7. FUTURE OBSERVATIONS

Future observations are required to see if indeed these stars continue to follow orbital arcs around each other proving that they are interacting with one another. In addition, future observations of the θ¹ Ori B₄ positions would help deduce if it is on a marginally stable orbit given its somewhat "non-hierarchical" location in the B group.

Future observations should also try to determine the radial velocities of these stars. Once radial velocities are known, one can calculate the full space velocities of these stars. Such observations will require both very high spatial and spectral resolutions. This might be possible with instruments such as the AO fed ARIES echelle instrument at the MMT.

We thank the whole Magellan Staff for making this wonderful telescope possible for use with our AO system. We especially thank Povilas Palunas for help over the entire MagAO commissioning run. Juan Gallardo (and his professional team), Patricio Jones, Emilio Cerda, Felipe Sanchez, Gabriel Martin, Maurico Navarrete, Jorge Bravo, and the whole team of technical experts who helped perform many exacting tasks in a very professional manner. Glenn Eychaner, David Osip and Frank Perez all gave expert support which was fantastic. We also thank Victor, Maurico, and Hugo for running the telescope so well. It is a privilege to be able to commission an AO system on such a fine telescope and site. Thanks from the whole MagAO team. We thank Miguel Roth, Francisco Figueroa, Roberto Bermudez, Sergio Veliz and, of course, Mark Philips for making this all happen—and very smoothly—despite the large AO team that was needed at the mountain and all the headaches and extra work it created for the LCO staff. We also thank the teams at the Steward Observatory Mirror Lab/CAAO (University of Arizona), Microgate (Italy), and ADS (Italy) for building such a great ASM. The MagAO ASM was developed with support from the NSF MRI program. The MagAO PWFS was developed with help from the NSF TSIP program and the Magellan partners. The Active Optics guider was developed by Carnegie Observatories with custom optics from the MagAO team. The VisAO camera and commissioning was supported with help from the NSF ATI program. L.M.C.'s research was supported by NSF AAG and NASA Origins of Solar Systems grants. J.R.M. is grateful for the generous support of the Phoenix ARCS Foundation. K.M. was supported under contract with the California Institute of Technology (Caltech) funded by NASA through the Sagan Fellowship Program.

DIFFRACTION-LIMITED VISIBLE LIGHT IMAGES OF ORION TRAPEZIUM CLUSTER WITH THE MAGELLAN ADAPTIVE SECONDARY ADAPTIVE OPTICS SYSTEM (MagAO)^{^*}

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ABSTRACT

1. INTRODUCTION

1.1. The Need for High-resolution Imaging