lbti: status and recent progress - university of...

Large Binocular Telescope Interferometer

LBTI: status and recent progress

D. Defrère and the LBTI team

University of Arizona

Tucson – March 23rd 2014


LBTI ComponentsLBTI Layout

18

Incoming Light

Visible

Light

Left Wavefront

Sensor

Incoming Light

Right Wavefront

Sensor

Visible Light

IR Light

NOMIC

(8-13 um)

LMIRCam

(3-5 um)

Phase Sensor

(2-2.4 um)

Beamcombiner

Nulling

Interferometer

2-2.4 and 8-13 um light

3-5 um light

IR Light

Nulling and Imaging

Camera (NIC)

Slow Corrector

(Piston, Tip-Tilt)

Fast (1 kHz )Corrector

(Piston, Tip-Tilt)

Imager


LMIRcam NOMIC

Wavelength Coverage (μm) 1.5-5.1 8-14 (8-25 capable)

Throughput >30% >20%

Pixel Size 0.011” 0.018”

FOV 20” 12”

Minimum Strehl 90% (3.8 µm) 98% (11 µm)

Spectral Resolution 350 100

5 sigma detection, 1 hour 19.0 (7 μJy) @ L’ 13.3 (200 μJy) @ N

Spatial Resolution 40 mas @ L’ 100 mas @ N’

LBTI science Cameras


LMIRcam NOMIC

Wavelength Coverage (μm) 1.5-5.1 8-14 (8-25 capable)

Throughput >30% >20%

Pixel Size 0.011” 0.018”

FOV 20” 12”

Minimum Strehl 90% (3.8 µm) 98% (11 µm)

Spectral Resolution 350 100

5 sigma detection, 1 hour ~2 MJ planet at 1 Gyr ~1 zodi debris disk

Spatial Resolution 0.4 AU at 10 pc 1 AU at 10 pc

LBTI science Cameras


First combined-aperture images from the LBT

Skemer et al. 2014


First coronagraphic image from the LBT

L’ PSF (DIT = 87ms)

0.7’’

L’ AGPM (DIT = 0.5s)

0.7’’


Infrared interferometric imaging of Io


Nulling interferometry at the LBTIFirst closed-loop null measurement in December 2013


First stabilized null measurements


1. LEECH: 100-night planet survey (see Skemer’s talk on Monday)

2. HOSTS: 60-night exozodi survey (this talk)

Skemer talk

Two large surveys


Why is NASA interested in exozodiacal dust?

• Source of noise and confusion for future exoEarth direct

imaging instruments:

1. Solar zodiacal cloud ~300 times brighter than Earth (IR and Visible);

2. Asymmetric features can mimic the planetary signal.

• LBTI is designed to probe for zodiacal dust down to the

level required to prepare future exoEarth imaging

instruments;

• Top level goal is to reduce the risk for future NASA

exoplanet imaging instruments;


Comparison of current facilities’ sensitivity to exozodiacal dust. LBTI’s

goal is to detect dust in the habitable zone down to 10 zodis.

Roberge et al.

NASA exoPAG report Dec. 2011

LBTI exozodi program in context

CHARA/PIONIER


Our first detection: h Crv (Feb. 2014)

tra

nsit

• Outer disk seen by Herschel (i = 46.8°, PA = 116.3°, Duchene et al.

2014);

• Excess: 17% (IRS), 4% (KIN);


Vega

TwilightHD 69830Clouds

b Umab Leo Twilight


Nuller performance status

• Demonstrated null depth of ~1% with a stability in the 0.3-

0.6% range;

• Null stability limited by ~11Hz vibration of secondary

swingarm.

Vega – Mar. 17 2014


(a) K And -- Bonnefoy et al. 2014, (b) HR 8799 --

Skemer et al. 2012, (c) NLTT 33370 -- Schleider et

al. 2013, (d) HK Tau b -- Leisenring et al. (in prep),

(e) HP Boo – Vaz et al. in prep, (f) VY Cma –

Shenoy et al. 2013, (g) Upper Sco – Bailey et al.

2013, (h) HD 32297 – Rodigas et al. 2014, (i) Eisner

et al. in prep., (j) HD114174 – Matthews et al. 2014

(see poster)

Environments of USco Brown Dwarf Companions 9

1 2 3 40

0.5

1

1.5

2

2.5

3

3.5x 10

−16

Wavelength [mm]

Fl [W/m2/

mm]

1700K

1800K

F igur e 6. 1RXS 1609B photomet ry and DUST Y SED comparisonwith A V = 0. T he green squares are the lit erature photomet ry andthe blue diamonds are this work. Spect ra are DUST Y 1800K and1700K model photospheres with log(g) = 4.0. T he filled circlesrepresent the integrated flux of each spect rum in that bandpass.

Tab le 4HIP 78530 system propert ies and apparent

photomet ry.

Property HIP 78530A HIP 78530B

Distance [pc] a 157± 13

Spect ral type b B9V M 8± 1

Mass b ⇠2.5 M 23± 3 M J u p

Tef f [K ] b ⇠10500 2800± 200

log(L / L ) b – -2.55± 0.13Separat ion [00] c 4.54± 0.09PA [◦ ] c 140.7± 1

J 6.928± 0.021 d 15.06± 0.05 b

H 6.946± 0.029 d 14.39± 0.04 b

K s 6.903± 0.020 d 14.17± 0.04 b

3.4 µm e 6.842± 0.034

L 0 c,f 6.87± 0.05 13.80± 0.064.5 µm g 6.88± 0.014.6 µm e 6.890± 0.0198.0 µm g 6.948± 0.008

12 µm d 6.924± 0.017

24 µm h 6.846± 0.01

a Hipparcos catalog (van Leeuwen 2007).b Lafreniere et al. (2011).c T his work.d 2M ASS J/ H / K S survey (Skrutskie et al. 2006).e System. W ISE survey (Wright et al. 2010).f MKO L 0.g System. Spitzer IRAC (Carpenter et al. 2006).h System. Spitzer M IPS (Carpenter et al. 2009).

L 0A = 6.87± 0.05. HIP 78530B was detected at a cont rast

between “ A” and “ B” of 6.93 magnitudes (Table 4). The5σ sensit ivity limit was L 0 = 16.21, giving SN R = 47and L 0

B = 13.80± 0.06. Our measured separat ion and PAare 4.5400± 0.0900 and 140.7◦ ± 1◦ , including plate scaleand distort ion uncertaint ies (see Sect ion 2.2.1), consis-tent with previous observat ions.

3.3.2. Spectral Classification and Near-IR Colors

The ext inct ion towards HIP 78530 is AV = 0.5 (Car-penter et al. 2009), and we dereddened the photometryusing this value. We calculated a K magnitude for the

F igur e 7. HIP 78530 LMIRCam L 0 image, smoothed for displaypurposes wit h a Gaussian kernel wit h FW HM = 1λ / D ⇠ 0.100. T hecompanion, circled, is 4.500 t o the SE.

BD of 14.12 ± 0.06 following the prescript ion describedin Sect ion 3.1.3.

Wefind thespect ral typesderived for HIP 78530B frombroadband 1− 4 µm colors and from near-IR spectra arenot consistent . HIP 78530B was classified using J/ H / Kspect ra as M8 ± 1 (Lafreniere et al. 2011), though theauthors noted the object was ⇠ 0.2 mag too blue at Kfor this spect ral type. We find that its H − L 0 color isalso ⇠ 0.5 mag too blue compared to an L2010 M8-typeBD (Figure 4, left panel). There is a similar mismatchin both H − K and K − L 0 colors (Figure 8, left panel).With theaddit ion of our new L 0 data, wefind theobject ’sJ through L 0 colors are not consistent with an M8, butinstead with an L2010 M3± 2. Int riguingly, the object ’sabsolute magnitude is consistent with other USco M8-type BD (Luhman & Mamajek 2012) as shown in Figure8 (right panel). We discuss the implicat ions of this dis-crepancy further in Sect ion 4.2. Regardless of whetherwe adopt spectral class M8 or M3, we find no evidencefor K or L 0 excesses.

3.3.3. 24 µm Flux

The observed and predicted fluxes for the HIP 78530system are 13.10 ± 0.12 mJ y and 12.0 ± 0.5 mJ y. Theexcess is 9± 4.5%. At 2σ, we do not consider it stat ist i-cally significant . This, combined with a lack of 1− 5 µmexcess around both “ A” and “ B,” suggests that neitherobject retains a massive warm disk.

3.3.4. Constraints on addit ional objects

No addit ional objects are detected in our data. Figure9 plots our new cont rast and mass limits as a funct ion ofprojected stellocent ric distance; we adopt the Hipparcossystem distance of 157 pc (van Leeuwen 2007). We reacha sensit ivity of 20 M J up beyond 100 AU and . 5 M J up

beyond 175 AU. In the east half of the system, our FOVextended to 800 AU, while on the west side, the FOVended at ⇠ 200 AU. An equally luminous binary com-panion to “ B” (i.e. “ Bb” ) is ruled out at > 15 AU sepa-rat ion from “ B.”

4. DISCUSSION

Characterizat ion of NLTT 33370 11

Fig. 1.— VLT -NaCo J (top-left ) and K s (bot tom-left ) and LBT -LM IRCam H (top-right ) and L 0-band (bot tom-right ) images of NLT T33370 A and B. T he separat ion of the components is⇢= 0.00076± 0.00005 in the J -band images. T he di↵erent epochs, separated by only twomonths, reveal the rapid orbital mot ion of the binary.

4

Fig. 2.— Top row: VY CM a imaged with LM IRCam for 2.9s t hrough the K s fil t er (λ0 = 2.15 µm), L filt er (λ0 = 3.8 µm) and M fi l t er(λ0 = 4.8 µm). Each image is a 10 × 10 FOV oriented North up, East left with linear greyscaling. T he bar in the lower left of eachimage measures 1 = 1200 AU at a distance of 1.2 kpc. T he beam size at L is 0. 12 (measured FW HM of a point source). T he dashedcircle in the center indicates where the image is saturated out to a radius of ∼ 0. 4 around the star. T he SW Clump appears bright in allt hree fi l t ers. T he patch of light appearing to the East of the star in each image is presumed to be from to an internal reflect ion ratherthan part of VY CM a’s ejecta, due to it s shift ing posit ions relat ive to the star between dit her posit ions and between fi lt ers. Bottom row:Images of SW Clump after subt ract ing a project ed 2-D profi le of t he star ’s flux into the region covering the Clump (see § 4.1).

the two methods does not change our finding that theClump’s flux at M can be largely accounted for by scat-tered light alone.

In the first method, we created an azimuthal averageof the radial profile of the central region containing thestar in each filter. The rangeof azimuth covering the SWClump (posit ion angles from 200◦ to 230◦ E of N) wasex-cluded from this average, as was the rangecontaining the“ ghost ” in the East of the images that is assumed to bedue to an internal reflect ion. We then rotated the profileof this azimuthal average to make a circularly symmetricimage to represent the star ’s light and subtracted it fromthe image.

In the second method, we masked out the region con-taining the SW Clump and replaced it with a 2-D surfacethat approximates the shape of the star’s 2-D profile intothe region covering theClump in each filter. The2-D starprofile surface was made using IDL ’sTRI SURF rout ine,which uses linear interpolat ion to create a surface fillingin the masked out region. We subtracted this projectedimage from the image in each filter. The subt racted im-ages are displayed in the bot tom row of Figure 2.

Figure 3 compares the two subt ract ion methods for acut through the SW Clump. At each wavelength, sub-t ract ing the 2-D projected star profile (red dashed line)results in a lower flux from the Clump than subt ract ingthe azimuthally averaged star profile (blue dashed line).At K s and L the 2-D projected star profile provides abet ter match to the likely profile of the star in the re-gion of the Clump, though it may slight ly oversubtractat radii close to the star. The difference in the flux of

the SW Clump at K s and L for the two methods is ∼20%. At M the azimuthally averaged star profile is sub-stant ially lower than the likely profile of the star in theregion of the Clump, which the 2-D projected star profilefollows more closely. The difference in the flux of the SWClump measured from the two methods is ∼ 50% at M.

While thesedifferencesmay be substant ial, they do notaffect our conclusion that the Clump’s brightness can beaccounted for largely with scat tered light alone. For thefollowing analysisof theSW Clump weadopt the averageflux from the two star subtract ion methods. In Figure4 we plot the flux of the SW Clump from the averageof the two subtract ion methods on VY CMa’s SED. Asa check on our calculat ion, we compared our flux valueswith those obtained by S01 from their lower resolut ionADONIS/ SHARP II+ ESO image of VY CMa at 2.14µm. Our flux-calibrated average brightness of the SWClump from the LMIRcam K s image agrees to within afactor of 2-3 with the surface brightness at the matchinglocat ion of the SW Clump in the ESO images.

4.2. Constraint on Thermal Flux of the SW Clump

LMIRCam’s K s, L and M filters sample the wave-length range where emission from dust grains may bedue to either scat tering, thermal emission, or a combina-t ion of both. Over most of the likely rangeof theClump’stemperature, we find that thermal emission contributesnegligibly to its flux in all three filters. This conclusionis driven by two factors: the very bright fluxes at K s andL which are due to scat tered light , and the far-IR fluxof the ent ire system, which constrains the thermal con-

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