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Gemini-LIGHTS: Herbig Ae/Be and Massive T Tauri Protoplanetary Disks Imaged with Gemini Planet Imager

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Published 2022 August 23 • © 2022. The Author(s). Published by the American Astronomical Society.
, , Citation Evan A. Rich et al 2022 AJ 164 109 DOI 10.3847/1538-3881/ac7be4

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earich@umich.edu

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1 Department of Astronomy, University of Michigan, West Hall, 1085 South University Avenue, Ann Arbor, MI 48109-1090, USA; earich@umich.edu

2 The University of North Carolina at Greensboro, USA

3 Astrophysics Group, University of Exeter, Stocker Road, Exeter, EX4 4QL, UK

4 Center for Astrophysics ∣ Harvard & Smithsonian, 60 Garden Street, Cambridge, MA 02138, USA

5 Physics Department, University of Michigan, Randall Lab, 450 Church Street, Ann Arbor, MI 48109-1090, USA

6 Department of Astronomy, University of Florida, Gainesville, FL 32611, USA

7 Department of Astronomy & Institute for Astrophysical Research, Boston University, 725 Commonwealth Avenue, Boston, MA 02215, USA

8 IPAC, Caltech, 1200 E California Boulevard, Pasadena, CA 91125, USA

9 Department of Physics and Astronomy, Rice University, 6100 Main Street, MS-108, Houston, TX 77005, USA

10 Leiden Observatory, Leiden University, 2300 RA Leiden, The Netherlands

11 Astrophysics Department, American Museum of Natural History, Central Park West at 79th Street, New York, NY 10024, USA

12 Departamento de Astronomía, Universidad de Chile, Camino El Observatorio 1515, Las Condes, Santiago, Chile

13 Department of Physics and Astronomy, University of Nevada, Las Vegas, 4505 S. Maryland Parkway, Box 454002, Las Vegas, NV 89154-4002, USA

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  1. Received 2022 April 8
  2. Revised 2022 May 26
  3. Accepted 2022 May 27
  4. Published

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Protoplanetary disks; Direct imaging; Exoplanets; Herbig Ae/Be stars; Circumstellar dust

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1538-3881/164/3/109

Abstract

We present the complete sample of protoplanetary disks from the Gemini- Large Imaging with the Gemini Planet Imager Herbig/T Tauri Survey, which observed bright Herbig Ae/Be stars and T Tauri stars in near-infrared polarized light to search for signatures of disk evolution and ongoing planet formation. The 44 targets were chosen based on their near- and mid-infrared colors, with roughly equal numbers of transitional, pre-transitional, and full disks. Our approach explicitly did not favor well-known, "famous" disks or those observed by the Atacama Large Millimeter/submillimeter Array, resulting in a less-biased sample suitable to probe the major stages of disk evolution during planet formation. Our optimized data reduction allowed polarized flux as low as 0.002% of the stellar light to be detected, and we report polarized scattered light around 80% of our targets. We detected point-like companions for 47% of the targets, including three brown dwarfs (two confirmed, one new), and a new super-Jupiter-mass candidate around V1295 Aql. We searched for correlations between the polarized flux and system parameters, finding a few clear trends: the presence of a companion drastically reduces the polarized flux levels, far-IR excess correlates with polarized flux for nonbinary systems, and systems hosting disks with ring structures have stellar masses <3 M⊙. Our sample also included four hot, dusty "FS CMa" systems, and we detected large-scale ( >100 au) scattered light around each, signs of extreme youth for these enigmatic systems. Science-ready images are publicly available through multiple distribution channels using a new FITS file standard that has been jointly developed with members of the Very Large Telescope Spectro-polarimetric High-contrast Exoplanet Research team.

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1. Introduction

Orbiting reservoirs of gas and dust are often observed around young stars and referred to as "protoplanetary" disks. As the name suggests, it has long been thought (e.g., Kant 1755) that planets form within such disks, though definitive evidence had been lacking until quite recently (PDS 70b,c: Keppler et al. 2018; Haffert et al. 2019). While the location for planet formation is not in doubt, the critical physical mechanisms at play are still hotly debated for a predictive theory of planet formation that can robustly explain the wild diversity seen in exoplanet demographics.

Protoplanetary disks were initially classified based on the shapes of their spectral energy distributions (SEDs). "Full disks" show a continuous spectrum resulting from thermal emission from ∼1500 K to 10 s of K. "Transition disks" contrast in the SED shape of "full disks" and lack near-infrared emissions due to a large gap or cavity close to the star (Strom et al. 1989; Calvet et al. 2002; Espaillat et al. 2014), possibly indicating planet formation or merely disk dissipation. With the advent of high-angular-resolution imaging by the Atacama Large Millimeter/submillimeter Array (ALMA) and with extreme adaptive optics systems on 8m class telescopes, a more complex picture emerges. Recent surveys such as the Disk Substructures at High Angular Resolution Project (DSHARP) and Disks ARound T Tauri Stars with SPHERE (DARTT-S) have revealed a host of substructure including spiral arms, rings, gaps, and nonazimuth asymmetries (Andrews et al. 2018; Avenhaus et al. 2018; Garufi et al. 2018). These features can be interpreted as signposts of forming planets; however exoplanets are not directly detected with the exception of PDS 70 b,c (Keppler et al. 2018).

Here, we bring the power of "polarized differential imaging (PDI)" to the study of planet formation. PDI reveals faint light scattered off the disk surfaces, cavity walls, and other dust structures in the circumstellar environment, revealing the three-dimensional distributions of small dust grains (Avenhaus et al. 2018; Rich et al. 2021). When coupled with ALMA imaging, which is sensitive to large dust grains settled into the disk midplane, PDI can monitor how dust grains grow and evolve with time and search for differences for systems with different stellar masses. Also, as PDI is dependent on illumination, the polarized flux will be influenced by shadowing of the inner disk onto the outer disk (Debes et al. 2017; Rich et al. 2019; Labdon et al. 2019; Muro-Arena et al. 2020) and the flaring angle of the outer protoplanetary disk.

In this work, we define the Gemini Large Imaging with the Gemini Planet Imager Herbig/T Tauri Survey (Gemini-LIGHTS) sample of 44 Herbig Ae/Be and T Tauri protostars imaged in near-infrared scattered light with the GPI (Section 2). Our new survey complements existing surveys (e.g., Strategic Explorations of Exoplanets and Disks with Subaru; Tamura 2009; DARTT-S, Avenhaus et al. 2018) by better populating the high-mass range (Herbig Ae/Be stars; >3 M⊙). We describe the reduction techniques utilized in our sample (Section 3) then present our calibrated images along with a descriptive analysis (Section 4). Next we limit our sample to targets with stellar masses between 1.4 and 8 M⊙ to search for trends between polarized flux and system characteristics (Section 5). Finally we discuss what trends we observe in the sample to help explain Herbig Ae/Be and T Tauri evolution (Section 6).

2. Gemini-LIGHTS Sample

The Gemini-LIGHTS sample was chosen to represent a broad range of T Tauri and Herbig Ae/Be stars with different protoplanetary disk structures including transition, pre-transition, and full disks. First, the sample consists of objects that were R <9 mag (GPI wave front sensor limit) and with decl. between +20° to −80° (lower airmass). Next, we identified these targets with significant infrared excess based on a color–color diagram of Wide-field Infrared Survey Explorer (WISE) and Two Micron All Sky Survey (2MASS) colors as shown in Figure 1. Additionally, we chose targets over a range of colors to achieve a mixture of transitional and full disks. Transitional disk targets host less near-infrared flux than far-infrared flux associated with a gap in the inner disk. This upper left portion of the diagram coincides with the location of the full disks shown in Figure 1. Full disk targets host equal near-infrared flux to far-infrared flux associated with no gaps in the disk. We note that a systems inclination can effect their broad SED categorization. Known equal-brightness binaries were not selected with a separation between 0farcs05 and 2farcs0 because they would inhibit the performance of the adaptive optics (AO) system. Compact binaries (e.g., HD 34700A) and unequal brightness binaries (e.g., FU Ori) were not selected against. Observations were taken in the J and H bands; however not every object has both bands observed. H-band observations were prioritized for dim R-band stars for better AO performance; otherwise J-band observations were prioritized.

Figure 1.

Figure 1. Sample parameters of the 44 targets in the Gemini-LIGHTS survey. (A) is an infrared color–color diagram using WISE and 2MASS colors. The dashed line represents a flat spectrum SED, and the gray shaded region are objects with no near-infrared or mid-infrared excess. (B) HR diagram with pre-main-sequence mass tracks (colored lines) and the zero-age main sequence (ZAMS) assuming solar-metallicity mass tracks from Bressan et al. (2012). Targets are classified as: stars with mass >8 M⊙ (orange x's), T Tauri stars (pink plus), and Herbig Ae/Be stars (green diamonds). (C) Histogram of estimated age. (D) Histogram of estimated stellar mass. (E) Age vs. stellar mass. (F) Histogram of target distances. Specific target parameter values can be found in Appendix A in Tables 2 and 3. Note that FU Ori is not plotted on the HR diagram as the Teff temperature is unknown.

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We also added archival observations of targets to our sample that have previously been observed by GPI. First we include an early science GPI project (Monnier et al. 2017), which includes data from MWC 275, HD 144432, MWC 863, and HD 169142. We have also included data found on the Gemini archive including targets PDS 66, HD 100546, HD 101412, HD 100453, HD 142527, and AK Sco, in which some have previously been published (Rodigas et al. 2014; Wagner et al. 2015; Follette et al. 2017).

The Gemini-LIGHTS sample includes 44 targets, which are summarized in Table 1, along with their stellar properties in Table 2, and photometry used for the sample selection in Table 3. Stellar properties are primarily taken from Vioque et al. (2018) to create a uniform sample. A Hertzprung–Russell (HR) diagram and histograms of the systems' properties are shown in Figure 1 where we also label which targets have masses >8 M⊙, Herbig Ae/Be stars (>7600 K), and T Tauri stars (<7600 K). Targets have distances between 40 and 5000 pc with the median distance of 350 pc, target age estimates span from 0.02 to 15 Myr and a median age of 2.3 Myr, and target stellar masses range from 0.3 to 20 M⊙ with a median value of 2.5 M⊙. The majority of our sample have masses consistent with Herbig Ae/Be stars. Also, all but FU Ori (∼0.3 M⊙) have central star masses larger than 1 M⊙.

Table 1. Gemini-LIGHTS Target Sample

TargetHD Name2MASS NameR.A.Decl.Distance
   (°)Decl. (°)(pc)
AK ScoHD 15240416544485-3653185253.6868708−36.8887086139.8 ± 0.6
CU ChaHD 9704811080329-7739174167.013825−77.6548583184.4 ± 0.8
FU Ori⋯05452235+090412386.34319869.0700691407.5 ± 3.0
GW OriHD 24413805290838+115212682.284962511.8701944408.0 ± 10.4
HD 34700 AHD 3470005194140+053842879.92253335.6452111350.5 ± 2.5
HD 36917HD 3691705344698-053414583.6957719−5.5707229450.4 ± 11.3
HD 37806HD 3780605410229-024300685.2595665−2.7168744401.6 ± 4.4
HD 38087HD 3808705430057-021845485.752413−2.3125911377.0 ± 5.4
HD 45677HD 4567706281742-130310997.0726019−13.0530934572.1 ± 14.6
HD 50138HD 5013806513340-0657592102.8891461−6.9664934351.0 ± 5.7
HD 85567HD 8556709502853-6058029147.6188579−60.96745511047.4 ± 18.0
HD 95881HD 9588111015764-7130484165.4900393−71.51341761109.9 ± 24.3
HD 98800HD 9880011220530-2446393170.5215893−24.777886442.1 ± 1.0
HD 98922HD 9892211223166-5322114170.631975−53.3698472650.9 ± 8.8
HD 100453HD10045311330559-5419285173.2730745−54.3246148103.8 ± 0.2
HD 100546HD10054611332542-7011412173.3558188−70.1947875108.1 ± 0.4
HD 101412HD 10141211394445-6010278174.9352141−60.1743876412.2 ± 2.5
HD 104237HD 10423712000511-7811346180.0211875−78.1929333106.6 ± 0.5
HD 139614HD 13961415404638-4229536235.1931708−42.4983306133.6 ± 0.5
HD 141569HD 14156915495775-0355162237.4905224−3.9213039111.6 ± 0.4
HD 142527HD 14252715564188-4219232239.1744971−42.323228159.3 ± 0.7
HD 142666HD 14266615564002-2201400239.1667625−22.0277806146.3 ± 0.5
HD 144432HD 14443216065795-2743094241.7414345−27.7195047154.8 ± 0.6
HD 145718HD 14571816131158-2229066243.2982588−22.4853188154.7 ± 0.5
HD 158643HD 15864317312497-2357453262.8540017−23.9627733125.7 ± 1.7
HD 169142HD 16914218242978-2946492276.1240709−29.7805404114.9 ± 0.4
HD 176386HD 17638619013892-3653264285.4122394−36.890856155.1 ± 0.7
Hen 3-1330HD 32682317065390-4236397256.7246196−42.61106311445.2 ± 40.0
Hen 3-225HD 7653408550867-4327596133.7862435−43.4666156885.0 ± 23.1
Hen 3-365HD 8764310043028-5839521151.1265−58.66452781589.1 ± 309.7
HR 5999HD 14466816083427-3906181242.142875−39.1050278158.6 ± 0.9
HT Lup⋯15451286-3417305236.3036208−34.2918389153.5 ± 1.3
MWC 147HD 25943106330519+101919998.271617410.3222059653.4 ± 11.6
MWC 166HD 5336707042551-1027156106.1063792−10.45437221219.6 ± 314.4
MWC 275HD16329617562128-2157218269.0886685−21.9562308101.0 ± 0.4
MWC 297⋯18273952-0349520276.9146958−3.831125417.8 ± 5.3
MWC 614HD 17921819111124+1547155287.796919315.7875669260.1 ± 2.2
MWC 789HD 25055006015998+163056790.499952716.5157295760.3 ± 28.5
MWC 863HD15019316401792-2353452250.0746503−23.8959517150.8 ± 0.5
PDS 66⋯13220753-6938121200.5314458−69.636719497.9 ± 0.1
Ty CrA⋯19014081-3652337285.4201238−36.8762242159.1 ± 4.4
V1295 AqlHD 19007320030250+0544166300.76046375.7379176847.9 ± 22.5
V921 Sco⋯16590677-4242083254.7782539−42.70234131482.4 ± 77.4
WRAY 15-535⋯10152198-5751427153.841516−57.86183814989.9 ± 478.1

Note. Distances from Gaia DR3 early release Gaia Collaboration et al. (2021).

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Table 2. Gemini-LIGHTS Target Characteristics

TargetMassAgeTefflog10(L)Disk Classification
 (M⊙)(Myr)(K)log10((L⊙)) 
AK Sco1.4${}_{-0.07}^{+0.07}$ (b)8.38${}_{-0.42}^{+1.72}$ (b)6250${}_{-250}^{+250}$ (b)0.62${}_{-0.01}^{+0.03}$ (b)Cont
CU Cha2.25${}_{-0.14}^{+0.11}$ (b)4.37${}_{-0.32}^{+1.11}$ (b)10500${}_{-500}^{+500}$ (b)1.54${}_{-0.06}^{+0.07}$ (b)Ring
FU Ori0.3 (f)2.0 (i)⋯⋯Irr
GW Ori2.8 (ab)5.0 (j)5500.0 (ag)1.68${}_{-0.08}^{+0.1}$ (ah)Irr
HD 1004531.25${}_{-0.06}^{+0.06}$ (b)6.53${}_{-0.49}^{+0.45}$ (b)7250${}_{-250}^{+250}$ (b)0.79${}_{-0.0}^{+0.02}$ (b)Spiral, Cont
HD 1005462.06${}_{-0.12}^{+0.1}$ (b)5.48${}_{-0.77}^{+1.41}$ (b)9750${}_{-500}^{+500}$ (b)1.37${}_{-0.05}^{+0.07}$ (b)Cont
HD 1014122.1${}_{-0.11}^{+0.11}$ (b)4.37${}_{-0.32}^{+0.22}$ (b)9750${}_{-250}^{+250}$ (b)1.58${}_{-0.04}^{+0.05}$ (b)Nondet
HD 1042371.85${}_{-0.09}^{+0.09}$ (b)5.48${}_{-0.4}^{+0.27}$ (b)8000${}_{-250}^{+250}$ (b)1.33${}_{-0.01}^{+0.04}$ (b)Undet
HD 1396141.48${}_{-0.07}^{+0.07}$ (b)14.49${}_{-3.6}^{+1.41}$ (b)7750${}_{-250}^{+250}$ (b)0.77${}_{-0.01}^{+0.03}$ (b)Cont
HD 1415691.86${}_{-0.09}^{+0.09}$ (b)8.62${}_{-1.19}^{+11.38}$ (b)9750${}_{-250}^{+250}$ (b)1.22${}_{-0.03}^{+0.03}$ (b)Ring
HD 1425271.61${}_{-0.08}^{+0.12}$ (b)6.63${}_{-1.55}^{+0.33}$ (b)6500${}_{-250}^{+250}$ (b)0.96${}_{-0.0}^{+0.03}$ (b)Spiral, Ring
HD 1426661.49${}_{-0.08}^{+0.08}$ (b)9.33${}_{-0.47}^{+0.77}$ (b)7500${}_{-250}^{+250}$ (b)0.94${}_{-0.05}^{+0.04}$ (b)Cont
HD 1444321.39${}_{-0.07}^{+0.07}$ (b)4.98${}_{-0.55}^{+0.25}$ (b)7500${}_{-250}^{+250}$ (b)0.97${}_{-0.01}^{+0.04}$ (b)Nondet
HD 1457181.6${}_{-0.08}^{+0.08}$ (b)9.8${}_{-0.49}^{+2.8}$ (b)8000${}_{-250}^{+250}$ (b)0.9${}_{-0.04}^{+0.05}$ (b)Cont
HD 1586433.35${}_{-0.22}^{+0.79}$ (b)1.22${}_{-0.57}^{+0.29}$ (b)9800${}_{-300}^{+900}$ (b)2.22${}_{-0.07}^{+0.26}$ (b)Nondet
HD 1691422.0${}_{-0.13}^{+0.13}$ (b)8.98${}_{-3.9}^{+11.02}$ (b)10700${}_{-900}^{+800}$ (b)1.31${}_{-0.22}^{+0.12}$ (b)Ring
HD 1763862.3${}_{-0.3}^{+0.14}$ (b)4.05${}_{-0.57}^{+15.95}$ (b)10700${}_{-900}^{+800}$ (b)1.58${}_{-0.22}^{+0.12}$ (b)Nondet
HD 34700 A2.66${}_{-0.13}^{+0.32}$ (b)1.4${}_{-0.44}^{+0.23}$ (b)5900${}_{-100}^{+110}$ (b)1.36${}_{-0.02}^{+0.1}$ (b)Spiral, Ring
HD 369173.71${}_{-0.75}^{+0.94}$ (b)0.99${}_{-0.5}^{+0.9}$ (b)11215${}_{-1316}^{+1109}$ (b)2.43${}_{-0.29}^{+0.24}$ (b)Nondet
HD 378063.11${}_{-0.33}^{+0.55}$ (b)1.56${}_{-0.6}^{+0.64}$ (b)10475${}_{-675}^{+1025}$ (b)2.17${}_{-0.14}^{+0.19}$ (b)Undet
HD 380873.21${}_{-0.38}^{+0.79}$ (b)1.75${}_{-0.64}^{+9.15}$ (b)13600${}_{-830}^{+2900}$ (b)2.19${}_{-0.22}^{+0.3}$ (b)Nondet
HD 456774.72${}_{-0.39}^{+1.19}$ (b)0.61${}_{-0.3}^{+3.77}$ (b)16500${}_{-750}^{+3000}$ (b)2.88${}_{-0.17}^{+0.32}$ (b)Cont
HD 501384.17${}_{-0.33}^{+0.46}$ (b)0.63${}_{-0.18}^{+0.19}$ (b)9450${}_{-450}^{+450}$ (b)2.46${}_{-0.09}^{+0.13}$ (b)Cont
HD 855676.32${}_{-0.39}^{+0.53}$ (b)0.22${}_{-0.05}^{+0.05}$ (b)13000${}_{-500}^{+500}$ (b)3.19${}_{-0.08}^{+0.1}$ (b)Undet
HD 958815.5${}_{-0.28}^{+0.5}$ (b)0.28${}_{-0.07}^{+0.05}$ (b)10000${}_{-250}^{+250}$ (b)2.85${}_{-0.07}^{+0.1}$ (b)Undet
HD 988000.7 (ab)8.5${}_{-1.5}^{+1.5}$ (h)4200 (al)0.33 (al)Undet
HD 989226.17${}_{-0.31}^{+0.37}$ (b)0.2${}_{-0.04}^{+0.01}$ (b)10500${}_{-250}^{+250}$ (b)3.03${}_{-0.05}^{+0.06}$ (b)Undet
Hen 3-13309.0⋯22280 (ai)5.3 (ai)Nondet
Hen 3-2257.46${}_{-0.37}^{+0.51}$ (b)0.17${}_{-0.03}^{+0.02}$ (b)19000${}_{-500}^{+500}$ (b)3.55${}_{-0.07}^{+0.09}$ (b)Nondet
Hen 3-36517.72${}_{-6.72}^{+10.87}$ (b)0.02${}_{-0.01}^{+0.05}$ (b)19500${}_{-3000}^{+5000}$ (b)4.6${}_{-0.53}^{+0.64}$ (b)Irr
HR 59992.43${}_{-0.12}^{+0.12}$ (b)2.73${}_{-0.35}^{+0.26}$ (b)8500${}_{-250}^{+250}$ (b)1.72${}_{-0.04}^{+0.05}$ (b)Nondet
HT Lup1.3${}_{-0.2}^{+0.2}$ (d)0.5${}_{-0.4}^{+0.02}$ (d)4247${}_{-237}^{+161}$ (k)0.51${}_{-0.01}^{+0.01}$ (k)Ring
MWC 1475.16${}_{-1.29}^{+1.84}$ (b)0.42${}_{-0.28}^{+0.53}$ (b)14000${}_{-2900}^{+2125}$ (b)2.97${}_{-0.4}^{+0.27}$ (b)Nondet
MWC 16612.3${}_{-4.2}^{+4.2}$ (e)0.08${}_{-0.08}^{+0.08}$ (e)29500${}_{-1000}^{+1000}$ (e)4.11${}_{-0.37}^{+0.37}$ (e)Nondet
MWC 2751.83${}_{-0.09}^{+0.09}$ (b)7.6${}_{-1.22}^{+1.05}$ (b)9250${}_{-250}^{+250}$ (b)1.2${}_{-0.03}^{+0.06}$ (b)Ring
MWC 29716.9${}_{-1.22}^{+1.87}$ (b)0.03${}_{-0.01}^{+0.01}$ (b)24500${}_{-1500}^{+1500}$ (b)4.59${}_{-0.12}^{+0.12}$ (b)Cont
MWC 6142.98${}_{-0.3}^{+0.18}$ (b)1.66${}_{-0.26}^{+0.54}$ (b)9500${}_{-200}^{+200}$ (b)2.05${}_{-0.14}^{+0.09}$ (b)Cont
MWC 7892.6${}_{-0.14}^{+0.3}$ (b)2.56${}_{-0.67}^{+0.43}$ (b)11000${}_{-500}^{+500}$ (b)1.94${}_{-0.12}^{+0.17}$ (b)Irr
MWC 8631.89${}_{-0.1}^{+0.1}$ (b)5.48${}_{-0.27}^{+0.44}$ (b)9000${}_{-250}^{+250}$ (b)1.37${}_{-0.04}^{+0.04}$ (b)Undet
PDS 661.2 (ak)6.0${}_{-1.0}^{+1.0}$ (g)5035${}_{-19}^{+19}$ (ak)0.0${}_{-0.01}^{+0.01}$ (ak)Ring
Ty CrA2.06${}_{-0.19}^{+0.22}$ (b)6.38${}_{-2.01}^{+13.62}$ (b)10700${}_{-900}^{+800}$ (b)1.41${}_{-0.23}^{+0.14}$ (b)Nondet
V1295 Aql5.89${}_{-0.76}^{+0.8}$ (b)0.22${}_{-0.07}^{+0.11}$ (b)9500${}_{-200}^{+200}$ (b)2.9${}_{-0.2}^{+0.16}$ (b)Nondet
V921 Sco19.96${}_{-5.0}^{+6.98}$ (b)0.02${}_{-0.01}^{+0.03}$ (b)29000${}_{-4500}^{+3882}$ (b)4.76${}_{-0.34}^{+0.33}$ (b)Undet
WRAY 15-53517.5${}_{-2.5}^{+2.5}$ (c)⋯20000${}_{-3000}^{+3000}$ (aj)6.0 (aj)Undet

Note. Stellar and disk characteristics of stellar mass, Age, Teff, Luminosity, and our disk classification. Disk classification described in Section 4.2. In the table Undet is short for Undetermined and Nondet is short for Nondetections. Citations for specific values are from the following: (b) Vioque et al. (2018), (c) Maravelias et al. (2018), (d) Garufi et al. (2020), (e) Fairlamb et al. (2015), (f) Zhu et al. (2007), (g) Murphy et al. (2013), (h) Ribas et al. (2018), (i) Beck & Aspin (2012), (j) Monnier et al. (2019), (k) Gaia Collaboration et al. (2018), (l) Takami et al. (2018).

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Table 3. Gemini-LIGHTS Photometry and Meeus Group

TargetWISE 2WISE 4 J band H bandMeeus Group
 (mag)(mag)(mag)(mag) 
AK Sco4.8 ± 0.0360.888 ± 0.0237.676 ± 0.0267.059 ± 0.033II
CU Cha4.3 ± 0.043−1.309 ± 0.0137.267 ± 0.0236.665 ± 0.049I
FU Ori3.509 ± 0.0650.175 ± 0.0216.519 ± 0.0235.699 ± 0.033⋯
GW Ori4.208 ± 0.045−0.74 ± 0.0117.698 ± 0.037.103 ± 0.029I
HD 1004533.388 ± 0.065−1.355 ± 0.0076.945 ± 0.0266.39 ± 0.038I
HD 1005463.156 ± 0.049−3.565 ± 0.0016.425 ± 0.025.962 ± 0.031I
HD 1014124.94 ± 0.0311.325 ± 0.0138.635 ± 0.0238.217 ± 0.047II
HD 1042372.469 ± 0.071−0.909 ± 0.0095.813 ± 0.0235.246 ± 0.059II
HD 1396145.099 ± 0.03−0.667 ± 0.0087.669 ± 0.0267.333 ± 0.04I
HD 1415696.469 ± 0.021.847 ± 0.0126.872 ± 0.0276.861 ± 0.04I
HD 1425273.066 ± 0.076−0.794 ± 0.0096.503 ± 0.0295.715 ± 0.031II
HD 1426664.145 ± 0.049−0.065 ± 0.0237.351 ± 0.0266.739 ± 0.027II
HD 1444324.398 ± 0.0440.042 ± 0.0167.095 ± 0.0326.538 ± 0.067II
HD 1457184.848 ± 0.040.585 ± 0.0157.69 ± 0.0247.263 ± 0.029II
HD 1586432.171 ± 0.091−0.028 ± 0.0144.9 ± 0.1864.712 ± 0.206II
HD 1691425.593 ± 0.023−0.561 ± 0.0087.31 ± 0.0216.911 ± 0.038I
HD 1763866.074 ± 0.015−0.455 ± 0.0126.847 ± 0.026.809 ± 0.031I
HD 34700 A6.794 ± 0.0160.993 ± 0.0088.041 ± 0.0237.706 ± 0.023I
HD 369174.278 ± 0.0441.153 ± 0.0257.221 ± 0.0196.964 ± 0.034II
HD 378063.419 ± 0.1150.116 ± 0.0517.115 ± 0.026.252 ± 0.033II
HD 380877.207 ± 0.0181.889 ± 0.0197.588 ± 0.0247.386 ± 0.042I
HD 456771.812 ± 0.03−2.895 ± 0.0017.242 ± 0.0266.347 ± 0.023I
HD 501381.585 ± 0.08−2.078 ± 0.0015.856 ± 0.0275.093 ± 0.029II
HD 855673.291 ± 0.0620.454 ± 0.0167.472 ± 0.0246.68 ± 0.031II
HD 958813.125 ± 0.0690.513 ± 0.0087.384 ± 0.0266.662 ± 0.044II
HD 988005.325 ± 0.0320.194 ± 0.016.397 ± 0.025.759 ± 0.027⋯
HD 989221.863 ± 0.014−1.054 ± 0.0066.004 ± 0.025.226 ± 0.029II
Hen 3-13303.344 ± 0.0631.95 ± 0.026.713 ± 0.0246.103 ± 0.042⋯
Hen 3-2257.066 ± 0.024.193 ± 0.0277.818 ± 0.0247.858 ± 0.04II
Hen 3-3652.248 ± 0.018−3.988 ± 0.0016.217 ± 0.0374.756 ± 0.268⋯
HR 59991.895 ± 0.012−0.541 ± 0.015.907 ± 0.0185.22 ± 0.027II
HT Lup4.973 ± 0.0320.856 ± 0.0167.573 ± 0.0216.866 ± 0.029⋯
MWC 1473.169 ± 0.074−0.757 ± 0.0137.454 ± 0.0266.666 ± 0.034II
MWC 1664.538 ± 0.0441.051 ± 0.0596.332 ± 0.026.22 ± 0.033⋯
MWC 2752.466 ± 0.061−0.758 ± 0.0076.195 ± 0.0215.531 ± 0.036II
MWC 2971.872 ± 0.018−4.718 ± 0.0016.127 ± 0.0194.387 ± 0.214⋯
MWC 6143.647 ± 0.051−1.606 ± 0.0066.994 ± 0.026.645 ± 0.026I
MWC 7894.633 ± 0.0390.046 ± 0.0158.475 ± 0.027.528 ± 0.026II
MWC 8633.244 ± 0.065−0.592 ± 0.0126.947 ± 0.026.214 ± 0.02II
PDS 666.183 ± 0.0221.587 ± 0.0168.277 ± 0.0327.641 ± 0.023⋯
Ty CrA5.139 ± 0.027−2.35 ± 0.0117.486 ± 0.0246.97 ± 0.026I
V1295 Aql3.443 ± 0.0640.589 ± 0.0167.194 ± 0.0196.647 ± 0.017II
V921 Sco2.045 ± 0.02−4.143 ± 0.0017.235 ± 0.0275.918 ± 0.045⋯
WRAY 15-5351.556 ± 0.011−0.231 ± 0.0115.762 ± 0.0184.959 ± 0.059⋯

Note. WISE photometry from the WISE All-sky survey (Cutri et al. 2012) and J- and H-band photometry from the 2MASS survey (Cutri et al. 2003). The Meeus Group is categorized in this work as described in Section 5. Notes: (a) Lieman-Sifry et al. (2016), (b) Henning et al. (1993), (c) Liu et al. (2018), (d) Fang et al. (2017), (e) this work, (f) Sylvester et al. (1996), (g) van der Plas et al. (2019), (h) Miley et al. (2019), (i) Nilsson et al. (2010), (j) Kataoka et al. (2016), (k) Ansdell et al. (2018), (l) Andrews et al. (2018), (m) Cazzoletti et al. (2019), (n) Benac et al. (2020).

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We note a limitation in our sample. We plot the estimated age versus stellar mass in Figure 1 and find that the two properties are highly correlated. This is likely due to our sample having few intermediate-mass T Tauri stars. This effect was noted by Guzmán-Díaz et al. (2020) where their recent survey of Herbig Ae/Be stars was also missing young intermediate-mass stars (see Figure 8 by Guzmán-Díaz et al. 2020).

2.1. Observations

Observations of the targets were taken at the Gemini South Observatory using the GPI (Macintosh et al. 2014; Poyneer et al. 2014; Larkin et al. 2014) in the J- and H-band filters. Observations were in the polarimetry mode using coronagraphic spots of sizes 184.7 mas and 246.7 mas for the J and H bands, respectively, and a pixel scale of 14.14 mas. Observations of the targets HD 100453 and HD 142527 were the only observations taken without a coronagraphic mask. Each image of a target measures the orthogonal polarization state of the 2farcs0 × 2farcs0 field of view (FOV). Between each image of a target, the half-wave plate is rotated 22fdg5. This creates a polarization set of four images at four different wave plate positions (0°, 22fdg5, 45°, 67fdg5). A typical observational epoch of a target observed eight sets of four polarization images producing 32 images per epoch. Exposure of the image for each target was adjusted for the targets brightness such that the PSF of the star does not saturate the images. Observations of the targets occurred between 2014 April and 2019 May with the majority of observations occurring between 2017 and 2019, and a list of the observations can be found in Table 4. Calibration files such as lamps and dark images were taken by the Gemini observatory staff every 2–4 weeks and accessed through the Gemini archive.

Table 4. Observational Log and Stellar/Instrumental Polarization

TargetFilterEpochExposure Time# of FramesStellar/Instrumental Polarization
  (YYYYMMDD)(s) % PolPA (°)
AK ScoH2018081159.6401.69 ± 0.06138 ± 1
CU ChaJ2017040658.2321.42 ± 0.12117 ± 3
CU ChaH2018041358.2361.5 ± 0.12112 ± 2
FU OriJ2018010358.2240.16 ± 0.0387 ± 6
GW OriJ2018010458.2320.48 ± 0.196 ± 6
GW OriH2018010458.2360.34 ± 0.11110 ± 10
HD 100453J2015041014.51400.07 ± 0.2241 ± 93
HD 100546J2017022043.6401.16 ± 0.0760 ± 2
HD 101412H2018031959.6400.68 ± 0.0592 ± 2
HD 104237J2017040758.2320.44 ± 0.0598 ± 3
HD 104237J2018031758.2320.31 ± 0.0194 ± 1
HD 104237J2018052058.2320.38 ± 0.04119 ± 3
HD 104237J2019021758.2320.51 ± 0.04121 ± 2
HD 139614J2017040658.2640.46 ± 0.09179 ± 5
HD 139614H2018060858.2280.56 ± 0.0311 ± 1
HD 139614J2019051358.2360.17 ± 0.08130 ± 14
HD 141569J2018060958.2320.65 ± 0.05105 ± 2
HD 141569J2019022358.2320.52 ± 0.06114 ± 3
HD 142527H201404254.4920.81 ± 0.1540 ± 5
HD 142666J2017070358.2200.49 ± 0.0684 ± 4
HD 142666J2018060958.2320.46 ± 0.0279 ± 1
HD 144432H2015070958.2321.26 ± 0.09175 ± 2
HD 145718J2018060758.2321.08 ± 0.16102 ± 4
HD 145718H2018060858.2321.51 ± 0.1882 ± 3
HD 158643J2018060958.2320.5 ± 0.061 ± 3
HD 158643J2019022358.2320.55 ± 0.023 ± 1
HD 169142J2014042558.2560.45 ± 0.13175 ± 8
HD 176386H2018060758.2320.28 ± 0.022 ± 2
HD 176386H2018081658.2320.27 ± 0.052 ± 5
HD 176386J2018081758.2640.15 ± 0.1440 ± 27
HD 34700 AJ2018010358.2320.65 ± 0.13102 ± 6
HD 34700 AH2018010358.2320.36 ± 0.05104 ± 4
HD 36917J2017123158.2320.21 ± 0.0331 ± 4
HD 36917H2019012658.2360.79 ± 0.0344 ± 1
HD 37806J2017123158.2320.89 ± 0.05129 ± 2
HD 38087J2018010358.2321.3 ± 0.06109 ± 1
HD 45677J2017123158.2320.42 ± 0.0446 ± 3
HD 45677H2018010158.2320.46 ± 0.0411 ± 3
HD 50138J2018010458.2320.29 ± 0.06143 ± 6
HD 85567J2018010358.2320.45 ± 0.0571 ± 3
HD 85567J2018052058.2320.46 ± 0.02128 ± 1
HD 85567H2019012758.2280.48 ± 0.0384 ± 2
HD 95881J2018010458.2400.91 ± 0.0986 ± 3
HD 98800H2019012734.9280.74 ± 0.1154 ± 4
HD 98922J2018032158.2320.41 ± 0.0764 ± 5
Hen 3-1330J2019051258.2322.02 ± 0.0627 ± 1
Hen 3-225J2019012758.2320.47 ± 0.05147 ± 3
Hen 3-225H2019012758.2320.21 ± 0.0839 ± 12
Hen 3-365J2017040658.2320.79 ± 0.09166 ± 3
HR 5999J2017040658.2321.13 ± 0.1128 ± 3
HR 5999J2017070258.2360.32 ± 0.1919 ± 17
HT LupH2019051452.4640.77 ± 0.1236 ± 4
MWC 147J2017123158.2320.79 ± 0.0288 ± 1
MWC 147H2019012758.2321.0 ± 0.0797 ± 2
MWC 166J2017040658.2160.81 ± 0.0548 ± 2
MWC 275J2014042458.2320.8 ± 0.3434 ± 12
MWC 297H2018060858.2321.21 ± 0.0594 ± 1
MWC 297H2018081743.6321.15 ± 0.0499 ± 1
MWC 614J2018081658.2640.94 ± 0.0795 ± 2
MWC 614H2018081658.2641.24 ± 0.11101 ± 2
MWC 789H2018112058.2400.92 ± 0.1337 ± 4
MWC 863J2014042458.2322.95 ± 0.2856 ± 3
PDS 66J2016030659.61040.82 ± 0.0699 ± 2
Ty CrAH2018060858.2320.33 ± 0.04148 ± 3
Ty CrAH2018081752.4320.15 ± 0.04164 ± 7
V1295 AqlH2018060858.2320.59 ± 0.06104 ± 3
V1295 AqlJ2018081658.2640.58 ± 0.0675 ± 3
V1295 AqlH2018081658.2320.56 ± 0.0453 ± 2
V921 ScoH2019051358.2320.98 ± 0.21138 ± 6
WRAY 15-535J2018031958.2322.1 ± 0.0953 ± 1
WRAY 15-535J2019022258.2322.02 ± 0.158 ± 1

Note. All epochs observed as part of the Gemini-LIGHTS survey. PA and %Pol are the average polarization angle and % polarization of stellar and instrumental polarization removed.

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3. Data Reduction

The data were primarily reduced using the IDL based Data Reduction Pipeline (DRP; Maire et al. 2010; Perrin et al. 2014) using version 1.5.0 (rev c0cad3f), written to reduce the GPI data. We also employed our own Python-based wrapper, which automated the execution of the DRP IDL pipeline to batch process our data to allow for consistent and reproducible reductions. The Python wrapper edits the DRP parameter files needed and executes the IDL commands. Changes to the standard reduction were based on previous work by Monnier et al. (2019) and Laws et al. (2020). The Python wrapper is available on Github. 14

In order to reduce our raw observational data, we inspected the individual raw frames and removed any frames if there was poor Adaptive Optics (AO) performance or if the target star slips to the edge of the coronographic mask. If a single frame was removed or missing from a polarization set, a frame from an adjacent set was used in its place. If more than one frame from a polarization set was removed or missing, the entire polarization set is not used. From raw images to polarization data cubes (PODC) files, we utilized the IDL DRP without modifications as described by Maire et al. (2010), Perrin et al. (2014, 2015), Millar-Blanchaer et al. (2016), and De Rosa et al. (2020). Each PODC file contains two images of the orthogonal polarization states. After PODC files were created, each observation was inspected by eye to verify if the flexure solution was correct.

The PODC cubes are centered using the IDL DRP code, which utilizes a Radon transformation of the satellite spots to calculate the location of the star behind the coronagraph. We found that this transformation could fail when there was a bright point source in the field near the satellite spots. We found that applying a circular mask of radius = 10 pixels over the point source was sufficient to correct for this issue. We tested the centering function on point sources in the FOV around HD 50138 and found a centroid accuracy of 0.22 pix (3.1 mas). This method was applied to all targets with a point source in the FOV. All centering solutions were inspected by eye to ensure the method worked properly. There were two instances where the Radon transformation within IDL that DRP failed to find the correct center with the companion masking improvements. The coronographic mask for HD 98800 was placed over the A component of the system rather than the disk hosting B component of the system (Kennedy et al. 2019). We centered the images on the disk hosting B component using an interactive Radon transformation algorithm described by Monnier et al. (2017). Second, the centering uncertainty for MWC 863 was large (1.9 pixels) using the IDL DRP Radon transformation. Using the interactive Radon transformation used from Monnier et al. (2017) improved the centering measured by the companions center (>1.2 pixels). We found the best solution was taking the average centering position from the IDL DRP Radon transformation and correcting for relative offsets of the PODC frames by measuring the center of the companion resulting in a centroid accuracy of (0.4 pixels). The relative centering method with the companion comes at the cost of greater uncertainty for the absolute position of the star behind the coronographic mask.

We next create Stokes cubes (I,Q,U,V) from the derotated (north-up) PODC files. Each set of half-wave plate positions (four PODC files) are double differenced into Stokes I, Q, and U, creating a set of approximately eight Stokes cubes per observing epoch. We note that while a circular polarization (V) image is created in the DRP pipeline, we do not use it in our analysis and will not refer to V further in this work.

We removed stellar and instrumental polarization (SIP) from the Q and U images. Measuring the amount of SIP is typically done by picking a region in the image and measuring the ratio of QSIP/I and USIP/I, which we will refer to as fQ and fU , respectively. The shape of the SIP should be similar to the I image, when dominated by the PSF; thus we can remove the SIP by multiplying the fraction of fQ and fU polarization to the intensity image and subtracting it from the Q and U images creating a set of corrected Q* and U*

Equation (1)

Equation (2)

where images have been corrected for SIP.

Next we picked a region in the image to best measure the SIP for each Q and U frame. As noted by Laws et al. (2020), using the region inside the coronagraphic mask region to estimate fQ and fU does not reliably remove the SIP for all targets. Laws et al. (2020) choose to use the region between 70 and 80 pixels away from the central star. This method is effective as long as the disk does not extend into this region or if there is a bright point source in this region. We chose to use a new method in which the entire FOV (0-140 pix) was used and mask out any regions where the Q/I or U/I ratio is larger than 0.05 or where there is known to be a point source within 10 pixels.

We tested if the masked method effectively measured fQ and fU by rotating the set of Stokes Q and U cubes into Qϕ and Uϕ frames where:

Equation (3)

Equation (4)

as defined by Monnier et al. (2019). This rotation will result in a distinctive quadrupole structure in Qϕ and Uϕ . If the masked method left a quadrupole structure, we utilized annulus regions of 10 pixels wide that minimized the quadrupole structure in the Qϕ and Uϕ images and recalculated Q* and U* described in Equations (1) and (2). These methods are compared in the appendix of Davies et al. (2022). Systems hosting bright point sources in the FOV occasionally exhibited some quadrupole structure in the Qϕ and Uϕ that could not be removed with the above methods (e.g., HD 98800 B and HD 144432). Future work is needed to better remove SIP from the GPI polarimetric data. For those wishing to replicate our reductions, the regions used to calculate fQ and fU can be found in the headers of our reduced FITS images along with the average fQ and fU removed. The averages for a given epoch fQ and fU are listed in Table 4 as the polarization angle (PA) and % polarization. We note that this information can be used to investigate the unresolved polarization of the inner disk region but caution that targets with low % polarization (<0.6%) will be dominated by instrumental polarization (Millar-Blanchaer et al. 2016).

Once the SIP is finally removed, the final Qϕ and Uϕ images are calculated, and the set of Qϕ and Uϕ are median combined to produce the final I, Qϕ , and Uϕ images as shown in Section 4. All I, Qϕ , and Uϕ images for every epoch can be found as a figure set appearing on the online version of this work (Figure 2). An example of the I, Qϕ , and Uϕ images can be seen in Figure 2. The fully reduced data can be found on Vizier and Data Behind the Figures (DBF). We also note that we have adopted a FITS header standard, in collaboration with the Very Large Telescope Spectro-polarimetric High-contrast Exoplanet Research (SPHERE) team members, which is presented in Appendix B to aid in easy comparison of data and result replication.

Figure 2.

Figure 2.

I, Qϕ , and Uϕ images for the Gemini-LIGHTS sample. Coronagraphic mask region marked with black circle. Location of likely binaries indicated in dashed black circles. Target and epoch date (YYYYMMDD) indicated in the upper left of each panel. The complete figure set (44 figures) is available in the online journal. (The data used to create this figure are available.) (The complete figure set (44 images) is available.)

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    Tarball of figure set images
    Figure 3.

    Figure 3. This figure shows the flux calibration scale for each of the targets and epochs in the sample. Blue dots are J-band observations, and orange dots are H-band observations. Observations that have their scale factors corrected (see Section 3.1) are noted with plus signs with their corrected values. Average J- and H-band scale factors are shown as dashed lines.

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    3.1. Flux Calibration

    The images are flux calibrated using the four satellite spots in each of the PODC images. To increase the signal to noise of satellite spots, the PODC files are mean combined. For H-band images, we used the first-order satellite spots. For J-band images, we used the second-order spots to ensure the spots were further away from the central PSF core. We note that the conversion factor between the satellite spots was updated since version 1.4 of DRP that was utilized by Laws et al. (2020) and Monnier et al. (2019). We used the known flux of the star from the 2MASS catalog to calculate the flux conversion factor and apply it to the Qϕ and Uϕ images, as shown in Figure 3. In our sample, we measure an average scale factor of 3.05 ± 0.57 for the J band and 3.09 ± 0.51 for the H band mJy ''−2/(ADU/s). We note that time (s) in the above scale factor is the total exposure time (ITIME×coadds).

    Central stars flux can be variable with time; thus without contemporaneous photometry we have taken effort to investigate and correct target epochs with discrepant flux calibration values. These different scale values could also be due to seeing conditions and AO performance. We investigated 14 observations that had a scale value 1σ (0.57 for the J band, 0.51 for the H band) away from the average scale value of 3.05 in the J band (3.09 in the H band). Four of these epochs had similar values to observations taken on the same night suggesting that the flux scale deviation is a correction for poor seeing or AO performance. Another four of these epochs had similar scale values to previous epochs suggesting that the divergent high or low scale value might be due to an over or underestimate of the flux of the target. The other six targets: FU Ori, HD 104237, HD 142666, HD 45677, HD 144432, and PDS 66 are likely to be due to variability; thus we use the average flux scale of other targets in that epoch. Finally, HD 100453 and HD 142527 were observed without coronographic spots; thus we choose to use flux calibration scale factors from Long et al. (2017) 0.836 mJy ''−2/(ADU/s) for both observations.

    3.2. Flat-field Accuracy

    Due to the design of the GPI instrument, flat-field corrections of the pixel-to-pixel variations in the raw frame are not currently possible. The current correction uses lamp flat observations to correct for low-frequency variations across the FOV. This issue was first investigated by Millar-Blanchaer et al. (2016). We observed the twilight sky to obtain sky flats in order to independently estimate the effects of the flat-field correction onto our data. We reduced the data using the same parameters as described above in Section 3. We plot an example of an individual Q and U frames as a % flux deviation from the average of the Q and U frames, as shown in Figure 4. We show that the low-frequency flat-field fails to correct for all large-scale flat-fielding variations, especially toward the edge of the detector. However, there are still sizable flux variations between 2%–4% at the center of the image. Thus we conclude that any azimuthal flux variations observed in the disk that are on the scale of 2%–4% may not be astrophysical but instrumental due to poor flat-fielding of the image.

    Figure 4.

    Figure 4. Showing the Q (top) and U (bottom) example sky flats taken on 2018 January 3. Flux plotted as the percent deviation from the flux average of the frame.

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    3.3. Uncertainty Propagation

    We estimate the uncertainty of our images by bootstrapping images from the fully reduced Stokes Qϕ and Uϕ images. We performed this bootstrapping 100 times and these sets of bootstrapped images are used for all analysis and error propagation in Section 4 and utilized in calculating the uncertainties of the polarization statistics for each observed epoch as shown in Table 5.

    Table 5. Polarization Flux and Point-source Detection Upper Limits

    TargetBandEpochQϕ /Fstar5σ Point-source Detection Upper limit
      (YYYYMMDD)(per 1000)Contrast at 0farcs2Mass Sensitivity ( M⊙)
    AK ScoH201808111.2 ± 0.21.6E-040.011
    CU ChaJ201704062.7 ± 0.93.9E-050.006
    CU ChaH201804133.8 ± 1.02.5E-040.02
    FU OriJ201801033.7 ± 0.61.2E-040.03
    GW OriJ201801049.4 ± 0.91.6E-040.011
    GW OriH2018010413.3 ± 1.15.9E-050.007
    HD 34700 AJ2018010315.7 ± 1.81.3E-040.008
    HD 34700 AH2018010322.2 ± 1.85.3E-050.005
    HD 36917J20171231<0.023.6E-050.007
    HD 36917H20190126<0.038.6E-050.01
    HD 37806J20171231<0.043.8E-050.008
    HD 38087J20180103<0.126.1E-050.008
    HD 45677J201712318.4 ± 1.15.5E-050.01
    HD 45677H201801016.9 ± 1.03.0E-050.01
    HD 50138J201801040.9 ± 0.26.0E-050.011
    HD 85567J201801030.10 ± 0.081.0E-040.03
    HD 85567J201805200.12 ± 0.074.7E-050.02
    HD 85567H201901270.13 ± 0.112.5E-040.06
    HD 95881J20180104<0.099.6E-050.03
    HD 98800H201901270.19 ± 0.14⋯⋯
    HD 98922J201803210.21 ± 0.166.9E-050.03
    HD 100453J201504102.9 ± 0.2⋯⋯
    HD 100546J2017022010.2 ± 0.68.7E-050.008
    HD 101412H20180319<0.041.7E-040.013
    HD 104237J201704070.04 ± 0.039.2E-050.013
    HD 104237J201803170.03 ± 0.021.6E-040.014
    HD 104237J201805200.06 ± 0.032.2E-040.02
    HD 104237J20190217<0.061.4E-040.014
    HD 139614J201704062.5 ± 0.78.9E-050.01
    HD 139614J201905131.5 ± 0.66.7E-050.009
    HD 139614H201806082.5 ± 0.73.1E-030.06
    HD 141569J201806090.29 ± 0.281.5E-040.013
    HD 141569J20190223<0.34.8E-050.007
    HD 142527H201404259.26 ± 1.17⋯⋯
    HD 142666J201707030.19 ± 0.17⋯⋯
    HD 142666J201806090.09 ± 0.08⋯⋯
    HD 144432H20150709<0.03⋯⋯
    HD 145718J201806070.4 ± 0.34.4E-050.007
    HD 145718H201806080.4 ± 0.21.1E-040.011
    HD 158643J20180609<0.012.4E-050.005
    HD 158643J20190223<0.04⋯⋯
    HD 169142J201404254.5 ± 0.81.0E-040.008
    HD 176386J20180817<0.012.8E-050.005
    HD 176386H20180607<0.056.7E-050.006
    HD 176386H20180816<0.026.8E-050.006
    Hen 3-1330J20190512<0.034.8E-050.04
    Hen 3-225J20190127<0.047.3E-050.013
    Hen 3-225H20190127<0.066.7E-050.01
    Hen 3-365J201704065.5 ± 0.63.3E-050.05
    HR 5999J20170406<0.023.3E-040.03
    HR 5999J201707020.03 ± 0.027.8E-050.01
    HT LupH201905140.4 ± 0.25.6E-040.009
    MWC 147J20171231<0.078.4E-050.011
    MWC 147H20190127<0.13.3E-040.04
    MWC 166J20170406<0.066.0E-031.4
    MWC 275J201404240.3 ± 0.27.3E-050.01
    MWC 297H201806080.3 ± 0.28.5E-050.04
    MWC 297H201808170.3 ± 0.21.0E-040.04
    MWC 614J201808162.0 ± 0.31.5E-040.014
    MWC 614H201808162.4 ± 0.21.7E-040.013
    MWC 789H201811205.9 ± 1.56.3E-040.06
    MWC 863J20140424<0.31.1E-030.03
    PDS 66J201603062.2 ± 1.15.9E-050.003
    Ty CrAH20180608<0.68.7E-040.3
    Ty CrAH20180817<0.441.5E-040.06
    V1295 AqlJ20180816<0.022.5E-040.05
    V1295 AqlH20180608<0.021.0E-040.03
    V1295 AqlH20180816<0.042.9E-040.05
    V921 ScoH201905130.06 ± 0.051.6E-040.13
    WRAY 15-535J201803190.11 ± 0.071.0E-04⋯
    WRAY 15-535J20190222<0.041.5E-05⋯

    Note. Mass estimate based on measured companion flux and the system age (see Table 2 and use models from Baraffe et al. (2015) for masses >0.01 M⊙ and Phillips et al. (2020) for masses <0.01 M⊙. The Qϕ /Fstar values utilized in this work are the weighted average of the epochs listed above.

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    3.4. Point-source Reduction

    We identified point sources in the FOV by looking for bright companions in the intensity image and identifying dimmer point sources using an angular differential imaging (ADI) reduction technique. For the ADI reduction, we used the Karhunen–Loéve image projection (KLIP) algorithm for epochs that had FOV rotation that was sufficiently large to take advantage of an ADI reduction. Using the centered PODC files from the standard reduction discussed above, we utilized the pyKLIP 15 package output (Wang et al. 2015). In order to avoid self-subtraction issues, we used parameters of subsection = 1 and annuli = 1, which minimize the number of individual segments in the FOV used to calculate the PSF. We found that 16 systems have point sources identified by the pyKLIP reduction. Upon close examination, all of these point sources are visible in the intensity images. A further 6 targets have point sources that were identified in the total intensity image but did not have sufficient FOV rotation for the pyKLIP analysis.

    We measured the flux and position of each point source by fitting a Gaussian to each point source in the intensity frame. The location of these point sources can be found in Qϕ and Uϕ images in the online figure sets (e.g., Figure 2) and are marked by dotted black circles. The point-source location, flux, and estimated mass are shown in Table 6. For the companion mass estimates we used models from Baraffe et al. (2015) for masses >0.01 M⊙ and the chemical equilibrium models from Phillips et al. (2020) for masses <0.01 M⊙.

    Table 6. Point Sources Detected in the Gemini-LIGHTS Sample

    TargetBandEpochComp. #SeparationPAFluxSaturatedMass Est.
      (YYYYMMDD) ('')(° )(mJy) ( M⊙)
    FU OriJ2018010310.431 ± 0.003163.3 ± 0.327 ± 4TRUE0.6 ${}_{-0.1}^{+0.1}$
    HD 100453J2015041110.909 ± 0.002131.63 ± 0.0915.4 ± 1.3FALSE0.08 ${}_{-0.04}^{+0.01}$
    HD 101412H2018031910.460 ± 0.002149.9 ± 0.21.46 ± 0.11FALSE0.053 ${}_{-0.003}^{+0.03}$
    HD 101412H2018031920.153 ± 0.003184.9 ± 1.05.0 ± 0.4FALSE0.19 ${}_{-0.01}^{+0.06}$
    HD 104237J2017040711.226 ± 0.004255.43 ± 0.1429 ± 9FALSE0.11 ${}_{-0.04}^{+0.04}$
    HD 104237J2018031711.214 ± 0.004255.034 ± 0.1319 ± 6FALSE0.065 ${}_{-0.02}^{+0.02}$
    HD 104237J2018052011.204 ± 0.003255.1 ± 0.1132 ± 6FALSE0.12 ${}_{-0.02}^{+0.02}$
    HD 104237J2019021711.198 ± 0.004255.28 ± 0.1134 ± 5FALSE0.13 ${}_{-0.02}^{+0.02}$
    HD 144432H2015070811.256 ± 0.0046.07 ± 0.1651 ± 7TRUE0.3 ${}_{-0.05}^{+0.03}$
    HD 158643J2018060910.118 ± 0.013294 ± 530 ± 4FALSE0.07 ${}_{-0.02}^{+0.01}$
    HD 38087J2018010310.303 ± 0.003183.7 ± 0.447 ± 8TRUE0.8 ${}_{-0.3}^{+0.5}$
    HD 50138J2018010410.729 ± 0.003106.6 ± 0.210.1 ± 1.5FALSE *0.12
    HD 98800H2019012710.374 ± 0.0109.8 ± 0.6633 ± 9TRUE0.4 ${}_{-0.03}^{+0.03}$
    HR 5999J2017040611.251 ± 0.004112.55 ± 0.11141 ± 16TRUE0.6 ${}_{-0.2}^{+0.1}$
    HR 5999J2017070211.260 ± 0.005112.4 ± 0.2128 ± 19TRUE0.6 ${}_{-0.2}^{+0.1}$
    HT LupH2019051410.161 ± 0.003246.6 ± 0.770 ± 7TRUE0.13 ${}_{-0.01}^{+0.01}$
    Hen 3-365J2017040610.818 ± 0.003101.08 ± 0.141.3 ± 0.2FALSE *0.17
    Hen 3-365J2017040620.446 ± 0.008133.0 ± 0.80.6 ± 0.1FALSE *0.13
    Hen 3-225J2019012711.564 ± 0.019303.4 ± 0.31.46 ± 0.02FALSE *0.12
    Hen 3-1330J2019051210.268 ± 0.01749 ± 20.8 ± 0.3FALSE *0.13
    MWC 147J2017123110.137 ± 0.00456 ± 130 ± 6TRUE0.6 ${}_{-0.1}^{+0.4}$
    MWC 147J2019012710.156 ± 0.00354.1 ± 0.833 ± 5TRUE0.5${}_{-0.1}^{+0.4}$
    MWC 166J2017040610.587 ± 0.004298.4 ± 0.3282 ± 9TRUE * > 1.4
    MWC 297H2018060810.573 ± 0.00386.5 ± 0.22.06 ± 0.09FALSE *0.04
    MWC 297H2018081710.576 ± 0.00486.6 ± 0.21.74 ± 0.14FALSE *0.04
    MWC 789H2018112010.371 ± 0.002216.8 ± 0.310.7 ± 0.8FALSE0.8 ${}_{-0.2}^{+0.1}$
    MWC 863J2014042110.975 ± 0.004226.97 ± 0.14154 ± 19TRUE0.8 ${}_{-0.1}^{+0.1}$
    Ty CrAH2018060810.129 ± 0.004255.2 ± 1.326 ± 4FALSE>1.4
    Ty CrAH2018081710.112 ± 0.003257.3 ± 1.041 ± 4FALSE>1.4
    V921 ScoH2019051310.475 ± 0.013319.0 ± 1.10.09 ± 0.03FALSE0.03 ${}_{-0.01}^{+0.01}$
    V921 ScoH2019051321.122 ± 0.003323.12 ± 0.090.96 ± 0.08FALSE0.15 ${}_{-0.01}^{+0.01}$
    V1295 AqlH2018060811.113 ± 0.00340.52 ± 0.110.026 ± 0.002FALSE *0.010
    V1295 AqlH2018081611.135 ± 0.01640.1 ± 0.50.055 ± 0.009FALSE *0.013
    V1295 AqlJ2018081711.119 ± 0.00540.6 ± 0.20.015 ± 0.002FALSE *0.010

    Note. Point sources with their PSF cores saturated are noted in the table. Mass estimate based on measured companion flux and the system age (see Table 2). Mass estimates use models from Baraffe et al. (2015) for masses >0.01 M⊙ and Phillips et al. (2020) for masses <0.01 M⊙. A * denotes any target with ages <0.5 Myr that are younger than the youngest model in Baraffe et al. (2015) or Phillips et al. (2020).

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    Using the pyKLIP routine, we also placed limits on the point-source companions we are sensitive to detecting in our sample. We calculated the 5σ flux limit of point sources at angular separations of 0farcs2 and 0farcs5 for all targets with sufficient field rotation. We define sufficient rotation where a point source would move >1 pixel due to field rotation, corresponding to a total field rotation of 4° at 0farcs2 and 1.6° at 0farcs5. We note that 6 systems, HD 98800 B, HD 100453, HD 142527, HD 142666, HD 158643, and WRAY 15-535, did not have sufficient field rotation at 0farcs2 to determine the contrast at a separation of 0farcs2. Based on the contrast and the systems age, we estimated the upper limit mass that would be detectable at 0farcs2 and 0farcs5. For angular separations of 0farcs2, in 9 systems we can detect down to Jupiter-mass companions, for 26 we can detect down to brown dwarf mass companions, and in 3 systems we can detect some stellar mass companions. For angular separations of 0farcs5, in all systems we can detect companions down to 1 Mjup mass companions. The results for each epoch can be found in the appendix in Table 5. For the companion mass estimate limits we used models from Baraffe et al. (2015) for masses >0.01 M⊙ and the chemical equilibrium models from Phillips et al. (2020) for masses <0.01 M⊙.

    4. Results

    We now present the complete sample of our results from the Gemini-LIGHTS survey. All I, Qϕ , and Uϕ images for all 71 epochs of the 44 targets can be found as a figure set appearing on the online version of this work (Figure 2). We will first identify targets in which we do not detect Qϕ flux. Next we will group the remaining targets by their disk morphological properties. We note that some targets can be found in multiple categories.

    Some of the observations from this long and large observational campaign have already been published as part of the Gemini-LIGHTS campaign. Laws et al. (2020) published observations of 4 large disks with irregular features (FU Ori, MWC 789, HD 45677, Hen 3-365). Monnier et al. (2019) investigated the J- and H-band observations of HD 34700A, which exhibit strong spiral structures. Davies et al. (2022) published the J- and H-band observations of HD 145718 that compared the scattered light images to photometry and near-IR interferometry and were constrained by radiative transfer modeling. Finally, Kraus et al. (2020) investigated the triple star system of GW Ori finding evidence of disk tearing.

    4.1. Nondetections

    Here we establish which targets have detected scattered light and which targets are nondetections. One advantage in rotating the polarization I, Q, and U into Qϕ and Uϕ is the ease of interpretation. The Qϕ flux should be dominated by photons that scattered only once and from one source, while the Uϕ flux could be dominated by photons that scattered multiple times, or scattering that is not azimuthally symmetric with the assumed central star (i.e., presence of a binary). We compute the amount of Qϕ and Uϕ flux within 0farcs4 (30 pixels) divided by the stellar flux for that given band. For sources with no close-in point source (e.g., within 0farcs4 of central star), for disks with very low Qϕ , the amount of Uϕ will go to zero before Qϕ , as shown in Figure 5. As Qϕ and Uϕ are drawn from the same observables, the noise in Qϕ is the same as Uϕ . Thus we look at the standard deviation of Uϕ for epochs <0.01% Qϕ and Uϕ and removed binaries by removing epochs >3σ. We then took the standard deviation of the summed Uϕ flux, which is also the uncertainty of the Qϕ flux. Using this metric, we establish that we do not detect polarized light around 18 of the 77 epochs using a 3σ deviation of Uϕ of 0.002% in the J band and 0.002% in the H band. Targets that are nondetections include: HD 36917, HD 101412, HD 144432, HD 158643, HD 176386, Hen 2-225, Hen 3-1330, MWC 147, MWC 166. Two targets, V1295 Aql (20180608 H band, 20180816 H band) and TY CrA (20180608 H band) where not detected in some epochs but were detected in other epochs (V1295 Aql: 20180816 J band; TY CrA: 20180817 H band). Thus, we do not detect polarization light around 9 of the 44 targets and additionally do not always detect polarization light in 2 of the targets.

    Figure 5.

    Figure 5. The % of Qϕ /I flux vs. Uϕ /I in a 0farcs4 annulus (30 pixels) for J-band (left) and H-band (right) observations. The dashed circle represents the 3σ significance.

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    4.2. Disk Structure Categorization

    We are categorizing our disks only by the observed disk structure in the Qϕ images. The categories include spiral armed disks, ringed disks, continuous disks, irregular disks, and undetermined disks structure. Our chosen categories are inspired by the disk structure categories defined by Garufi et al. (2018). We have deviated when appropriate to better describe our sample. The results of our classification are tabulate in the appendix in Table 2.

    4.2.1. Spiral Armed Disks

    We find there are four targets that host one or more spiral arms including HD 100453, HD 139614, HD 34700 A, and HD 142527, as shown in Figure 6. HD 100453 hosts symmetric spiral arms, while HD 34700A and HD 142527 most multiple arms that are not symmetric. HD 139614 has an arm structure on one side of the disk. All four of these systems have had the origin of their spiral arms investigated previously (HD 100453: Wagner et al. 2015; HD 139614: Laws et al. 2020; HD 34700 A: Monnier et al. 2019; A. S. E. Laws 2022, in preparation; HD 142527: Long et al. 2017).

    Figure 6.

    Figure 6. Spiraled Disks: Qϕ images of disks classified as hosting spiral arms including targets HD 100453, HD 139614, HD 34700 A, and HD 142527. See online figure sets (e.g., Figure 2 for target specific color bars). The image flux scale is log and chosen to highlight the target structure. The point sources in the FOV are labeled with a dashed black circle. The name of the target and the epoch in YYYYMMDD format can be found in the upper left of each image. The type of polarized image (Qϕ ), the scale of the image (in au), and the photometry band (J or H) can be found along the bottom of each subimage.

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    4.2.2. Ringed Disks

    We define disks that have one or more concentric polarized light dust rings in their Qϕ image, as shown in Figure 7. We note that we deviate from the definition of Garufi et al. (2018) as they make a distinction between ringed and rimmed systems; however scattered light imaging lacks the information to robustly make the distinction due to the inner working angle. We find that seven targets show signatures of rings: HD 169142, HD 141569, MWC 275, CU Cha, HD 34700 A, HD 142527, and PDS 66. We note that two of the four spiral armed disks are also ringed disks (HD 142527, HD 34700 A).

    Figure 7.

    Figure 7. Ringed Disks: Qϕ images of disks classified as hosting rings including targets HD 169142, HD 141569, MWC 275, CU Cha, HD 34700 A, HD 142527, PDS 66. See online figure sets (e.g., Figure 2 for target specific color bars). See Figure 6 caption for figure details.

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    4.2.3. Continuous Disks

    We define continuous disks as disks that are resolved and do not appear to have a gap or hole in their PDI or show a strong nonazimuthal structure, as shown in Figure 8. Some of these disks may have a gap or ringed structure, but it is not shown in the polarized light imagery due to the inclination (HD 145718: Davies et al. 2022) or inner working angle (IWA; MWC 614: Kluska et al. 2018). We note that this does not imply that these disks do not have gaps at all, but those gaps are not visible in our polarized light imaging. We find that 11 targets show signatures of rings: AK Sco, HD 45677, HD 50138, HD 100453, HD 100546, HD 139614, HD 142666, HD 145718, HT Lup, MWC 297, and MWC 614. We note that one of the spiral armed disks are also continuous disks (HD 139614).

    Figure 8.

    Figure 8. Continuous Disks: Qϕ images of disks classified as continuous including targets: AK Sco, HD 45677, HD 50138, HD 100453, HD 100546, HD 139614, HD 142666, HD 145718, HT Lup, MWC 297, and MWC 614. See online figure sets (e.g., Figure 2 for target specific color bars). See Figure 6 caption for figure details.

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    4.2.4. Irregulars

    We define irregular disks as disks that host structures that have strong nonazimuthal features and disks that are very large in size (>300 au). There are four disks in our sample that exhibit these features: MWC 789, FU Ori, GW Ori, and Hen 3-365, which are shown in Figure 9. GW Ori irregularity is thought to be due to disk tearing as investigated by Kraus et al. (2020).

    Figure 9.

    Figure 9. Irregular Disks: Qϕ images of disks classified including targets: MWC 789, FU Ori, GW Ori, and Hen 3-365. See online figure sets (e.g., Figure 2 for target specific color bars). See Figure 6 caption for figure details.

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    4.2.5. Undetermined

    There are some targets that have a significant amount of polarized light detected (see Section 4.1); however the polarized light in the image does not extend far enough away to reliably categorize the disk into one of the above categories. Objects that fit this description include: HD 37806, HD38087, HD 85567, HD 95881, HD 98922, HD98800, HD 104237, HR 5999, MWC 863, TY CrA, V921 Sco, V1295 Aql, and WRAY 15-535, as shown in Figure 10. We note this is a similar to the "small disk" category utilized by Garufi et al. (2018). However, our sample is not strictly distance limited; thus some objects have disks that do appear to be small (e.g., HD 104237 disk radius <21 au), while other objects such as V921 Sco appear to be distant and have rather large disks (<300 au). Thus we choose to classify these objects as undetermined.

    Figure 10.

    Figure 10. Undetermined disk structures: Qϕ images of disks that are insufficiently resolved to determine their morphology. Targets include: HD 37806, HD38087, HD 85567, HD 95881, HD 98922, HD98800, HD 104237, HR 5999, MWC 863, TY CrA, V921 Sco, V1295 Aql, and WRAY 15-535. See online figure sets (e.g., Figure 2 for target specific color bars). See Figure 6 caption for figure details.

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    4.2.6. Categorization Summary

    Utilizing the polarized disk morphology described above, we now investigate if we see any trends with the disk morphology. We plotted all 44 targets on an HR diagram and an infrared color–color diagram in Figure 11. We note that targets can be in multiple categories. The first trend is that ringed systems only occur in systems with masses <3 M⊙. This could partially be due to a distance effect where the ringed structure is within the IWA of the observations, though less likely as there are six systems with resolved disks that do not exhibit any ringed structures. Second, irregular systems appear to be younger; however, this is not a definitive trend as only three objects on the HR diagram are classified as irregular.

    Figure 11.

    Figure 11. HR diagram (left) and color–color diagram (right) of this work's entire sample. Each target is marked by its apparent polarized light disk structure described in Section 4.2. HR diagram includes pre-main-sequence mass tracks (colored lines) and the zero-age main sequence (ZAMS) assuming solar-metallicity mass tracks from Bressan et al. (2012). Note that FU Ori is not plotted on the HR diagram as the Teff temperature is unknown.

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    Next, we plot the complete sample on a color–color diagram as shown in Figure 11. We generally find that full disks are unresolved, undetected, or continuous disks. Pre-transitional disks are mostly continuous disks with some irregulars and ring disks. Finally, transitional disks are dominated by ringed disks, as expected.

    4.3. FS CMa stars

    Four of our targets have been identified as FS CMa targets, which are a subtype of B[e] stars: HD 45677, HD 50138, HD 85567, and HD 98922 (Miroshnichenko et al. 2007; Vioque et al. 2020). These stars are potentially post-main-sequence stars. We note that HD 45677 is also known as FS CMa, the prototype for this classification. We have the first resolved images of the dust around HD 45677, HD 50138, HD 85567, and HD 98922, as shown in Figure 12. All objects have significant Qϕ flux that is detected out to sizable distances of ∼250 au (HD 50138), ∼700 au (HD 45677), ∼300 au (HD 85567), and ∼180 au (HD 98922). Additionally, each of the objects also show significant Uϕ flux, which is suggestive of multiple scatterings due to optical depth effects or multiple illumination sources due to a binary. The Uϕ flux for both HD 45677 and HD 50138 is likely to be due to optical depth effects as the pattern is within the location of high Qϕ flux similarly seen in other optically thick sources such as HD 34700 A, GW Ori, and HD 100546. HD 50138 does have a point source located within 1farcs0 of the central star but is not co-located with the concentric Qϕ flux around the central star. There is excess polarized flux north of the main disk that could be faint additional structure (e.g., back of the disk, spiral arm), but we are unable to definitively determine its origin. HD 85567 is very distant (1047 pc); thus only the very outer portion of the system is imaged. Finally, HD 98922 exhibits a centro-symmetric Qϕ pattern around the inner working angle.

    Figure 12.

    Figure 12. Qϕ and Uϕ images of targets that have previously been classified as FS CMa. The image scale is log and chosen to show target structure. See the figure set appearing on the online version of this work (Figure 2) for target specific color bars. The point sources in the FOV are labeled with a dashed black circle in Qϕ images. The name of the target and the epoch in YYYYMMDD format can be found in the upper left of each image. The type of polarized image (Qϕ or Uϕ ), the 1" scale and the size in au, and the photometry band (J or H) can be found along the bottom of each subimage. In Uϕ images, red is positive and blue is negative flux.

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    4.4. Point-source Detections in the Field of View

    We found 24 point sources within the FOV of the PDI images with 21 targets hosting these point sources. Point sources are identified in Qϕ and Uϕ images in the online figure sets (e.g., Figure 2) with a dotted black circle marking the companions locations. The values are also tabulated in Table 6 located in the appendix. We assessed the likelihood that a given point source was a background star from the measured separation and brightness from the target star. We measured the local star density by using the number of stars per half magnitude within 1 deg of the target from the 2MASS catalog. We then calculated the probability that a background star would be located within the projected distance of the target. We found that 22 of the 24 point sources are likely not background stars within 3σ. Two point sources found around V921 Sco are less than a 2σ probability not being background stars. Based on the likelihood that all of the point-source companions are gravitationally bound; for the purposes of this work we will assume that all of the 24 point sources are binaries. However, future follow-up observations, especially V921 Sco, are needed.

    We estimated the masses of each of the companions using the age of the system and the companions measured flux and extrapolated the mass from evolutionary models from Baraffe et al. (2015) for masses >0.01 M⊙ and Phillips et al. (2020) for masses <0.01 M⊙. We estimated the mass uncertainties using the uncertainty in the systems age and flux. We note that some of our targets have ages less than 0.5 Myr, which is younger than the youngest age by Baraffe et al. (2015). In these cases, we estimated their mass based on the youngest age in the Baraffe et al. (2015) models to give an idea of their mass but note these masses should not be used for future analysis as they are not reliable. Estimated masses are tabulated in Table 6 located in the appendix.

    We compare the estimated companion mass and the projected companion separation to the stellar mass, as shown in Figure 13. Of our 24 point sources, one point source around V1295 Aql has a mass consistent with a super Jupiter, four have masses consistent with brown dwarf masses (HD 101412, HD 158643, MWC 297, and V921 Sco), 12 systems have companions with masses consistent with M dwarfs (0.08–0.57 M⊙), and seven have masses >0.57 M⊙. Further, we note that more massive stellar objects (>6 M⊙) have larger projected separations due to these objects being more distant. For these objects with separations >1000 au, future observations are necessary to determine if these objects are comoving and if they are gravitationally bound.

    Figure 13.

    Figure 13. Comparison of the stellar mass and estimated companion mass (left) and of the stellar mass and projected companion separation (right). For the companion masses, we denote the boundary between brown dwarfs and stars (hydrogen burning limit: blue) and denote the boundary between planets and brown dwarfs (deuterium burning limit: orange).

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    4.5. Possible Nonconcentric Reflection Nebulae

    Two targets, MWC 147 and TY CrA, have polarized flux that is adjacent to the target, as shown in Figure 14. The flux plotted is flux that is 1σ per pixel above the noise in the image. Both targets are known to be associated with reflection nebulae; thus the flux is likely material that is part of the larger reflection nebulae. We also note that the flux may be an instrumental effect due to known issues with flat-fielding (see Section 3.2). However, this explanation is unlikely as the features appear in the same location over multiple epochs and multiple bands, and no other sources show these features taken on the same night.

    Figure 14.

    Figure 14. Target adjacent polarized flux of TY CrA (top row) and MWC 147 (bottom row) for two different epochs. Both targets have point sources (dashed circles) close to the inner working angle (solid circle). The companion and flux <1σ are masked in the images. The flux is in mJy "−2.

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    4.6. Systems with Stellar Masses >8 M⊙

    Six of our systems have stellar masses >8 M⊙ including Hen 3-365, Hen 3-1330, MWC 166, MWC 297, V921 Sco, and WRAY 15-535. We do not detect polarized light around two of the six systems (Hen 3-1330 and MWC 166), and we classified the morphology of the rest as one irregular (Hen 3-365), one continuous (MWC 297), and two undetermined (V921 Sco, WRAY 15-535). Several of these objects have conflicting classifications between young stellar objects (YSOs) and post -main-sequence stars. Hen 3-365 is discussed by Laws et al. (2020) who noted that it is inconsistent with being a massive supergiant due to the large H I column density, but parallax distances are consistent with being an evolved star (Oudmaijer et al. 1998; Maravelias et al. 2018). Hen 3-1330 has previously been classified as a WR+O binary system (Richardson et al. 2011). V921 Sco has been classified as both a supergiant and as a Herbig Be star (Kreplin et al. 2020). WRAY 15-535 has been classified as a supergiant B[e] star (Domiciano de Souza et al. 2007). Finally, previous work investigating MWC 166 and MWC 297 has shown that they are consistent with being Herbig Ae/Be stars (Manoj et al. 2007; Wichittanakom et al. 2020). We leave the polarized light investigations of these individual targets to future work.

    4.7. Scattered Light Diagnostics

    We calculate the total amount of Qϕ flux within the image relative to the stellar flux Fstar. To mitigate noise being added into the images, we sum all of the Qϕ light between the IWA of the target and where the radial profile of the Qϕ is within 3σ of zero flux using error propagation from the bootstrapped images discussed in Section 3.3. The summed Qϕ flux is then divided by the flux of the star using the 2MASS J- and H-band flux. The Qϕ /Fstar ratio can be found in the appendix in Table 5. We note that this simplified methodology does not account for the r2 flux loss. However, this method can easily be applied to our entire sample and is not dependent on knowing the disk geometry (i.e., disk inclination), which is not possible for some systems (i.e., unclassified, irregular; see Section 4.2). As there are multiple epochs of some targets, we will use the weighted average of Qϕ /Fstar ratio for all analysis.

    We first compare the Qϕ /Fstar ratio to the color–color diagram shown in Figure 15, which we utilized for our target selection described in Section 2. The size of the circle corresponds to the Qϕ /Fstar ratio on a logarithmic scale. We find that the targets with the largest Qϕ /Fstar ratio are commonly found in the middle of the color–color plot between 3< W2–W4 <6 mag and 2 < J–W2 <4 mag. This region coincides with the transitional disk region as outlined in Figure 1. Low and nondetected Qϕ /Fstar ratios are scattered throughout the diagram but have a higher concentration at the bottom with W2–W4 <3 mag. This lower-right portion of the diagram coincides with the location of the full disks shown in Figure 1. Interpretation of these results will be discussed in Section 6.

    Figure 15.

    Figure 15. Infrared color–color plot (left) of the targets where the size of the circle corresponds to the magnitude of the Qϕ /Fstar ratio. Note that the size of the circles is logarithmic. The color–color diagram is the same as that shown in Figure 1 where colors are using 2MASS J-band magnitudes and WISE 2 (4.6 μm) and 4 (22 μm) bands representing near-infrared excess (J–W2) and mid-infrared excess (W2–W4). The dashed line represents a flat spectrum SED. The gray shaded region marks objects with no near-infrared or mid-infrared excess. The color–color diagram (right) labels each target with the corresponding number and target name on the bottom of the figure. The Qϕ /Fstar ratios plotted are the weighted average for each target. The individual Qϕ /Fstar values can be found in Table 2.

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    5. Scattered Light Trends

    We next compare the total amount of summed scattered light per stellar flux, Qϕ /Fstar, to disk and stellar parameters to find which parameters affect the amount of scattered light flux. Due to the diverse nature of our sample, there are some very-high-mass and very-low-mass stars. Here, we restrict our sample to only include stars between A-type stars (>1.4 M⊙) up to stars that will produce supernovae (<8 M⊙). When we have this restriction, we have a sample size of 33 targets that are also all in the Vioque et al. (2018) sample, which provides a common source of temperatures, ages, and masses. In this section we will only analyze these 33 targets.

    We will compare our subsample of 33 targets to their binarity, defined as either an outer binary where the stellar/substellar companion can be directly imaged in our sample (separation >0farcs12) or an inner binary (separation <0farcs12). Binarity is determined using high-contrast imaging (see Section 3.4) or is previously determined in the literature using spectroscopy, imaging, or interferometry. For the 33 targets in this sample, we are able to detect binaries with masses down to brown dwarfs (0.075 M⊙) at a separation of 0farcs2. There are three exceptions, HD 142527, HD 142666, and HD 158643, which did not have sufficient field rotation to determine contrast upper limits at 0farcs2. However, given the depth and that all of our binary targets identifiable in the intensity frame without the aid of ADI, we expect a similar sensitivity for these targets as well. We note that the statistics of inner binaries is not complete as we are dependent on previous studies, and it can be notoriously difficult to detect Doppler shift of binaries in YSOs. Additionally, our analysis is not complete for outer binaries with separations larger than the FOV (separation >1farcs9) and would not be detected by our survey.

    Lastly, we investigate the commonly used Meeus Group I and Group II categorization (Meeus et al. 2001) of Herbig Ae/Be SEDs. Group I disks have much larger far-infrared fluxes than Group II disks, and this classification is a version of the W2–W4 color but with only two categories. Here, we calibrated our WISE-based Group I and Group II classification (see Table 2) using the previous classification by Guzmán-Díaz et al. (2020), which has 29 of our targets in their Herbig Ae/Be sample.

    5.1. Infrared Colors

    We first compare the Qϕ /Fstar ratio to the W2–W4 color shown in Figure 16. We find that bluer targets are more likely to have lower Qϕ /Fstar ratios than redder targets. By our group definition above, Group II targets are more likely to have lower Qϕ /Fstar ratios, while Group I targets are more likely to have higher Qϕ /Fstar ratios. While there is a positive trend between the Qϕ /Fstar ratio and the W2–W4 color, there is a large spread of values in the trend especially noting the logarithmic plots in Figure 16. We attribute this large spread due to the presence of binaries. When the binaries are removed from the sample, as shown in the upper right in Figure 16, there is not as significant of a spread in values. This suggests that binarity complicates the correlation between polarized flux from the disk versus infrared color. Very bright binaries can make it difficult to remove the SIP, as discussed in Section 3, possibly resulting in a larger spread in Qϕ /Fstar ratio values. However, this is unlikely to be the main cause of the spread as the majority of binaries in these systems are not bright enough to leave the quadrupole residual in Qϕ and Uϕ .

    Figure 16.

    Figure 16. 33 targets with 1.4 M⊙< stellar Mass <8 M⊙ comparing the Qϕ /Fstar ratio to the system parameters of WISE 2 (W2, 4.6 μm)–WISE 4 (W4, 22 μm) color (upper right), W2–W4 color of only nonbinary systems (upper left), system age/ZAMS age (bottom left), and stellar mass (bottom right). The target shapes and colors refer to their classification (Group I-blue squares, Group II-orange circles, FS CMa-green diamonds) and the shape filling refers to the binarity (filled shapes: inner and outer binary; edge filled: outer binary; inner filled: inner binary; no filling: no binary). Outer binaries are exterior to the IWA and in the FOV as discussed in Section 3.4. Inner binaries are spectroscopic binaries from the literature. Error bars are plotted under the symbols with down pointing errors indicating the 2σ upper limit uncertainty.

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    There are two clear exceptions to this trend, TY CrA and HD 176386. Both of these targets have no detected Qϕ /Fstar ratios; yet both are very red (W2–W4 >6 mag). TY CrA and HD 176386 are both located in an extended reflection nebulae. Images from unWISE of TY CrA exhibit extended emission offset from TY CrA (Lang et al. 2016). Also, HD 176386 is most likely blended with CrA-54 in WISE bands, where CrA-54 is a K7 star that is known to host a disk (Cazzoletti et al. 2019). Thus we conclude that the W2–W4 color for TY CrA and HD 176386 are likely contaminated from their surrounding environment (reflection nebulae, close-by stars, dust clumps), not originating from the protoplanetary disk dust around these stars. However, we lack observations of adequate spatial resolution to provide accurate WISE-band fluxes; thus we will keep these targets in our analysis but label them as suspicious WISE fluxes.

    5.2. Age

    We next compare the summed Qϕ /Fstar flux to the ages of the systems. Stars will evolve toward the zero-age main sequence (ZAMS) at different rates with 8 M⊙ stars evolving much quicker than 1.4 M⊙ stars. Thus we plot the ratio of the current age of the stars divided by the target ZAMS for a given estimated stellar mass, as shown at the bottom left of Figure 16. For reference, the ZAMS is also plotted on an HR diagram for the sample in Figure 11. There are no systems with Qϕ /Fstar ratios <0.2 that are on the main sequence (>1.0 Age/ZAMS). We do not see a correlation between the system age/ZAMS and the Qϕ /Fstar ratio.

    5.3. Stellar Mass

    Finally, we compare the stellar mass to the ratio of Qϕ /Fstar shown at the bottom right of Figure 16. We find that the less massive stars have the largest Qϕ /Fstar. The one key result is looking at systems with stellar masses between 4 < Mstar < 8 M⊙, where only five systems have significantly detected Qϕ /Fstar. Four of these five systems are also classified as FS CMa stars. We will discuss this point further in Section 6.6.

    6. Discussion

    6.1. Ringed Systems and Stellar Mass Dependence

    We found that, in our sample, protoplanetary disks hosting ringed systems only occur in systems with stellar masses <3 M⊙ (see Section 4). There are several potential causes that could explain this trend. First, we note that dust ringed systems have largely been theorized to exist due to the formation of exoplanets creating gaps in-between these dusty rings. Recent evidence for this theory has been provided in the case of PDS 70b,c (Keppler et al. 2018). One explanation is that the physical scenarios that make it possible for exoplanets to form visible rings at lower stellar masses (<3 M⊙) do not exist for higher-stellar-mass systems. One possibility is the increased binarity for stellar masses >3 M⊙ disrupts the formation of rings and gaps, which was similarly observed for millimeter dust grains as part of the DSHARP sample (Kurtovic et al. 2018). However, in our sample we see the same fraction of binaries around more massive stars (42% for 3 M⊙ < stellar mass <8 M⊙) than we do around lower-mass stars (40% for Stellar Mass <3 M⊙) though we are not complete to inner and outer binaries as noted in Section 5. Second, the temperature of the higher-mass disks may prevent them from forming rings and gaps. Disks around higher-mass stars may be hotter, which can prevent the most volatile ice species, like CO, from freezing out onto dust grains, as found recently in a comparison of carbon depletion in T Tauri versus Herbig AeBe disks (van der Marel et al. 2021; Sturm et al. 2022). If the dust rings are created by the accumulation of millimeter-sized grains at the edges of gaps, which may in turn be carved by planets, reduced amounts of CO or CO2 ice in the warmer disks could inhibit planetesimal formation or grain growth. This could inhibit outer disk ring formation for the higher-mass disks. Garufi et al. (2018) found that polarized light rings only appeared around older stars (>5 Myr); thus what we may be observing is that stellar systems with masses >3 M⊙ do not have the time to develop polarized light rings before the dust has been removed from the system. Also, theoretical predictions suggest that the mass of a planet needed to open a gap in a disk is proportional to the mass of the star (see Equation (1) in Matsumura & Pudritz 2003). Thus the lack of gaps could be representative of more massive planets being needed to open up gaps around more massive stars.

    Finally, the lack of detected ringed systems for stellar masses <3 M⊙ could be an observational bias. In our sample, more massive targets are more likely to be more distant. Thus these more distant objects may host ringed structures but are located inside the IWA of our observations. Of the 20 systems with stellar masses >3 M⊙ (see Figure 11), all rings in the two systems with stellar masses <3 M⊙ could be detected, and an additional 11 systems are close enough that at least one ring could be detected. Our sample has a dust ring occurrence rate of 30% for systems with <3 M⊙ stars; thus we would expect to detect at least one system with stellar mass >3 M⊙ hosting rings. However, the detectability of rings around more massive systems is complicated because the central stars are also more luminous, increasing the outer disk brightness. Therefore ringed structures around systems with masses <3 M⊙ that are too dim to be detected by our imaging surveys, such as those found around MWC 275 (330 au; Rich et al. 2019) and CU Cha (341 au; Ginski et al. 2016), would be brighter and detectable around more luminous systems. If we assume the distance of the furthest known ring (341 au), we could detect such a ring around 15 of 20 systems with stellar masses >3 M⊙. Our survey is the first survey to search for trends in protoplanetary disks with stellar masses >3 M⊙; thus future work needs to verify these findings and search for massive systems that host ringed and gaped structures.

    6.2. Polarized Flux and Infrared Color

    We find a correlation between the polarized flux and the infrared colors, as shown for the entire sample in Figure 15. This trend replicates previous studies by Garufi et al. (2017, 2020), which found a similar correlation. For the entire sample, we find that objects associated with full disks are more likely to have a lower Qϕ /Fstar ratio, as shown in Figure 15. The lower Qϕ /Fstar ratio is expected as full disks are thought to self-shadow, causing the outer portions of the disk imaged with GPI to be dimmer in polarized light as compared to transitional disks where self-shadowing is not occurring.

    A similar trend can be seen for our limited Herbig Ae/Be sample discussed in Section 5. The strong infrared flux of Group I objects has often been interpreted as due to strong disk flaring but could also be due to disk cavity, in either case leading to easily detected scattered light disk flux (Maaskant et al. 2013; Garufi et al. 2017, 2018). The bluer Group II objects can be explained either by strong self-shadowing by the inner disk or a flatter geometry due to dust growth/settling (e.g., Muro-Arena et al. 2018). This would broadly correspond to larger scattered light flux from Group I objects and less from Group II objects. Our observations shown in Figure 16 broadly replicate this trend though there is a large spread in values. This is primarily due to the presence of binaries. Further, we observe very few binaries that are in redder objects (Group I for Herbig's) and most binaries in our sample are associated with bluer objects, as shown in Figure 16. This strongly suggests that a system hosting a binary plays a significant role in the Group I versus Group II classification. One potential cause is that the presence of binaries truncate the outer disks causing systems with binaries to have less polarized fluxes or not be resolved. This would match with previous studies suggesting that Group I and Group II systems are two distinct evolutionary pathways (Maaskant et al. 2013; Garufi et al. 2018).

    6.3. Polarized Flux and Age

    We observe a lack of systems with ages in the main sequence (>1.0 Age/ZAMS) and low Qϕ /Fstar ratios, as discussed in Section 5. This could be an indicator that systems with low Qϕ /Fstar ratios and older ages (>1.0 Age/ZAMS) do not have sufficient infrared excess and are thus not in our sample. However, the brightest systems with large Qϕ /Fstar ratios may have long-lived disks. This is similar to the conclusion found by Garufi et al. (2018), where the most massive disks are long-lived and easily observed in polarized light to older ages. Next, we do not find the correlation between Qϕ /Fstar and age, which had previously been observed by Garufi et al. (2018). However, our sample contains a larger range in stellar mass of 0.3–20 M⊙, whereas the Garufi et al. (2018) study only contained targets with stellar masses <3 M⊙. Thus, the correlation between polarized light and age may not be present for higher-mass stars. Finally, age results are difficult to interpret for more massive stars due to the lack of intermediate-mass T Tauri stars in our sample (see Section 2). Future work is needed to include intermediate-mass T Tauri stars in polarized light imaging surveys to assess age trends in Herbig Ae/Be systems.

    6.4. Polarized Flux and Binarity

    We find that only 5 of the 12 systems with stellar masses between 4 M⊙ < Mstar < 8 M⊙ have significantly detected Qϕ /Fstar ratios. There are four possible explanations for this trend. First, our more massive stars tend to be further away; thus, we may not be resolving as much of the polarization flux compared to closer systems. Second, these more massive stars tend to have more outer binaries (5 of 12) than less massive systems (<4 M⊙). These binaries could be truncating or stripping the disks from the central star. The lack of polarized light may be caused by stellar evolution, where the quickly evolving bright star photoevaporates the disks so they are no longer visible. Finally, the systems may host large disks that our imaging would resolve. However, no polarized light is detected because the inner disk shadows the outer disk. This mechanism has previously been invoked to explain nongaped protoplanetary disks having little polarized light detected (Garufi et al. 2017; Muro-Arena et al. 2018).

    6.5. Point Sources in Literature

    We compare our point-source findings to those found in the literature. We limit our discussion to the substellar mass companions and notable stellar companions. Our survey has imaged five substellar point sources (V921 Sco, HD 158643, MWC 297, HD 101412, and V1295 Aql). We found two companions around V921 Sco at 0farcs5 (0.03 M⊙) and 1farcs1 (0.15 M⊙). The first companion with a mass consistent with a brown dwarf at 0farcs5 was previously discovered by Ubeira-Gabellini et al. (2019) with a similar mass estimate (0.06 M⊙) but has not previously been confirmed until this work. They did not observe the second companion at 1farcs1; however this is expected as the second companion is located outside of the FOV of Ubeira-Gabellini et al. (2019) images. The brown dwarf object (0.07 M⊙) HD 158643 has not been directly imaged previously. However, Kervella et al. (2019) measured a variation in the proper motion of the system consistent with an object with a mass of 0.05 M⊙ normalized to a separation of 1 au. These values are potentially consistent, but further analysis is necessary. Ubeira-Gabellini et al. (2020) previously confirmed that MWC 297 B was comoving and measured the companion mass to be ${0.25}_{-0.15}^{+0.25}$ M⊙. This is much larger than our estimated mass of 0.03 M⊙. However, they were also able to measure a local high extinction of AV ∼ 11.9 mag from the spectral slope of the planet, which is sufficient to match our findings. Finally, there is no known previous discovery of the planetary mass companion around V1295 Aql or a brown dwarf companion around HD 101412. Follow-up observations to confirm that these objects are comoving is necessary.

    Ty CrA is a known quadrupal system with Chauvin et al. (2003) observing the fourth stellar companion around Ty CrA. The Chauvin et al. (2003) work was unable to verify if the object was comoving. Given the 16 year difference between our observations and Chauvin et al. (2003) observations and assuming the proper motion of Ty CrA, we find that if the fourth component was not comoving, the object would have moved by 0farcs52, while we measure a positional difference of 0farcs21. Given the large timescale and that the fourth component of Ty CrA is at a projected separation of ∼18 au, it is likely we are observing the orbital motion of the fourth component of Ty CrA. However, confirmation of the object being bound is necessary given the systems complex orbital dynamics.

    6.6. Classification of FS CMa stars

    For the four systems that are FS CMa candidates in our sample, HD 45677, HD 50138, HD 85567, and HD 98922, we observe polarized flux at large projected distances (>200 au), as shown in Figure 12. Additionally, these polarized signals are bright, and in the case of HD 45677 and HD 50138 there is significant Uϕ polarized flux suggesting multiple scattering events due to the optically thick nature of the material.

    Previous studies have detected extended structures around B[e] stars via H-alpha (Marston & McCollum 2008; Liimets et al. 2022). These structures are typically shell-like or indicative of bipolar outflows. We do not observe similar structures around our four FS CMa candidates, which appear to be closer to typical protoplanetary disk structures. However, previous studies were at a larger spatial scale (>1') and emission, while our observations are at smaller spatial scales and from a polarized source. Finally, our targets have a lower mass than those studied by Marston & McCollum (2008) and Liimets et al. (2022).

    Stars that exhibit the B[e] phenomenon are a heterogeneous group that include pre-main-sequence, main-sequence, and evolved systems. Additionally, many of these stars exhibit similar features making classification difficult. FS CMa as defined by Miroshnichenko et al. (2007) hosts emission-line spectra containing hydrogen lines; large infrared excess that peaks at 10–30 μm, located outside of a star-forming region; and if it has a secondary companion it is either fainter and cooler than the primary or degenerate. However, even given these parameters, distinguishing between B[e] classifications can be difficult. For example, a thorough investigation by Varga et al. (2019), which studied HD 50138, concluded that the evolutionary state of HD 50138 could not be unambiguously determined through mid-infrared spectroscopy. We conclude that the structures observed around the FS CMa candidate stars in our Qϕ and Uϕ images are more similar to the expected structures of protoplanetary disks rather than outflow from FS CMa or another type of evolved star. Future investigations into the resolved circumstellar material around stars are necessary including resolved ALMA observations to test the kinematics of the gas in the systems. The classification of these systems to be protoplanetary is extremely important as they represent four of the five systems with 4 M⊙ < Mstar <8 M⊙ in our sample that have detected scattered light flux.

    7. Summary

    We have presented the Gemini-LIGHTS survey, which observed 44 bright Herbig Ae/Be stars and T Tauri stars with the GPI instrument at Gemini South. We constructed our sample based on their near- and mid-infrared colors selecting a number of transitional, pre-transitional, and full disks in order to create an unbiased sample. Importantly, we did not select against unequal mass binaries with moderate separations. Observations of these 44 targets were taken in J and H bands utilizing high-contrast polarized imagery. Our selection criteria facilitate an unbiased approach to studying bright Herbig Ae/Be and T Tauri stars that does not favor famous targets with large disks (>100 au). We discussed how we uniformly reduced and analyzed the 70 epochs of observations that are part of the sample using the GPI DRP pipeline along with our own Python wrapper. We discussed improvements to finding centers of images with bright companions and improvements in removing SIP.

    Within our sample of 44 targets, we found several significant trends:

    • 1.  
      We detect scattered light signatures around 80% of our 44 targets.
    • 2.  
      Systems with large separation binaries are more likely to have bluer mid-infrared colors (Group II), and these systems are more likely to have no detectable Qϕ /Fstar fluxes.
    • 3.  
      We find that all ringed classified systems have stellar masses <3 M⊙, potentially indicating that if the disk rings arise from planet formation, the planet formation process or disk evolution is different around more massive stars.
    • 4.  
      We find a similar infrared (WISE 2–WISE 4) correlation with Qϕ /Fstar identified by Garufi et al. (2017). We find a large spread in values primarily due to binary systems as the trend is much tighter when binaries are removed.
    • 5.  
      Four of five of our targets with 4 M⊙ < Mstar < 8 M⊙ for which we have detected scattered light flux are also classified as FS CMa (HD 45677, HD 50138, HD 85567, and HD 98922). Due to the large radial extent of the polarized flux in the images (>200 au), we conclude that these objects are likely young systems.
    • 6.  
      We detect 24 point sources consistent with being star companions. We find 1 point source around V1295 Aql that is consistent with a super-Jupiter mass and 3 point sources consistent with brown dwarf masses. We confirm the existence of two brown dwarf candidates (V921 Sco, HD 158643) from previous direct imaging and proper motion discoveries. We find one brown dwarf candidate HD 101412, which had previously not been observed before.

    Thanks to Christian Ginski for their collaboration on defining the FITS header standard for high-contrast scattered light imaging. We would also like to thank Bruce Macintosh, Fredrik Rantakyro, Marshall Perrin, Max Millar-Blanchaer, Tom Esposito, Robert De Rosa, Jeffrey Chilcote, and René Rutten for their help on this survey. E.A.R. and J.D.M. acknowledges support from NSF AST 1830728. A.A. acknowledges support from NSF AST-1311698. S.K. acknowledges support from an ERC Consolidator Grant (Grant Agreement ID 101003096). L.P. gratefully acknowledges support by the ANID BASAL projects ACE210002 and FB210003, and by ANID—Millennium Science Initiative Program—NCN19_171. We thank the anonymous referee for feedback that helped to improve this paper. This work has made use of data from the European Space Agency (ESA) mission Gaia (https://www.cosmos.esa.int/gaia), processed by the Gaia Data Processing and Analysis Consortium (DPAC; https://www.cosmos.esa.int/web/gaia/dpac/consortium). Funding for the DPAC has been provided by national institutions, in particular the institutions participating in the Gaia Multilateral Agreement.

    This work is based on observations obtained at the international Gemini Observatory, a program of NSFs NOIRLab, which is managed by the Association of Universities for Research in Astronomy (AURA) under a cooperative agreement with the National Science Foundation on behalf of the Gemini Observatory partnership: the National Science Foundation (United States), National Research Council (Canada), Agencia Nacional de Investigación y Desarrollo (Chile), Ministerio de Ciencia, Tecnología e Innovación (Argentina), Ministério da Ciência, Tecnologia, Inovações e Comunicações (Brazil), and Korea Astronomy and Space Science Institute (Republic of Korea). List of program ID's where the data were obtained: GS-2017A-LP-12, GS-2017B-LP-12, GS-2018A-LP-12, GS-2018B-LP-12, GS-2015A-C-1, GS-2016B-DD-5, GS-2015B-Q-501, GS-2014B-Q-503, GS-2018A-FT-101-17, GS-2018A-FT-101-15, GS-2018A-FT-101-5, GS-2014A-SV-414-14, GS-2014A-SV-414-8, GS-2014A-SV-406-5, GS-2017B-Q-500-16, GS-2014A-SV-412, and GS-2015A-Q-49.

    Appendix A: Target Properties

    In this section, we provide the target parameters, observation log, and calculated values used in this work. Table 1 lists all of the targets, HD names, their coordinates, and their GAIA distances. Table 2 lists the targets of the Gemini-LIGHTS sample along with target system properties such as the stellar mass, age, effective temperature, and luminosity. The photometry used in this work is listed in Table 3.

    Table 4 lists each of the 70 observational epochs observed as part of the Gemini-LIGHTs sample. The table lists each target, filter, date of observation, length of image exposure, and number of frames. Additionally, the table lists the average % polarization that was removed during the reduction. Note that this value is a combination of SIP and for small values of % polarization, the polarization will be dominated by the instrumental polarization.

    Table 5 lists the measured Qϕ /Fstar ratio for every observed epoch. A description of how Qϕ /Fstar is calculated can be found in Section 5. Additionally, the 5σ upper limit of potential point sources detected at a separation of 0farcs2 is listed in Table 5. Details on the ADI reduction and analysis can be found in Section 3.4. Table 6 lists the point sources detected in the Gemini-LIGHTS sample.

    Appendix B: Polarized Light FITS header Definition

    Here we define a standard FITS file for high-contrast polarization data. This is motivated by large numbers of observations of protoplanetary disks from multiple telescopes (e.g., GPI, SPHERE/IRDIS, Subaru/HiCIAO, Subaru/CHARIS). A standard FITS file will allow for better comparison of protoplanetary disk polarization imagery between different studies and instruments. This FITS standard was created in collaboration with Christian Ginski.

    The data are held in a three-dimensional cube of 5 × x × y, where x and y are the pixel dimensions of the image, and 5 are the different image types. The five different image types are I, Qϕ , Uϕ , Q, U, and :LP_I where I is the intensity image without stellar light subtracted, Qϕ , Uϕ , Q, and U as defined above in Section 3, and LP_I, which is the linear polarized intensity or ${({Q}^{2}+{U}^{2})}^{1/2}$.

    A sample header is shown in Table 7 taken from GPI observations of MWC 275 (HD 163296) used in this work. Standard WCS headers are included to allow use and image overlay with programs such as DS9. Additionally, reduction information such as star locations (X-STAR, Y-STAR), flux calibration (ZEROPT, REFMAG, CALIBFFAC), and fraction of stellar/instrumental polarization removed (FQ, FU) will help facilitate better comparisons between different epochs of observations of the same target.

    Table 7. Sample Header

    KeywordValueComment
    DATE-OBS"2014-04-24"UT start date of exposure
    CCDSIZE"2048x2048"Array dimensions
    CREATOR"GPI DRP, v1.5.0,revc0cad3f"This file created by GPI Data Reduction
    OBSMODE"J_coron"Currently selected observation mode
    RA269.0887Target R.A.
    DEC−21.956075Target decl.
    DATE"2014-04-24"UTC Date of observation (YYYY-MM-DD)
    EPOCH2000.0Target Coordinate Epoch
    GEMPRGID"GS-2014A-SV-412"Gemini program ID
    INSTRUME"GPI"Instrument used to acquire data
    OBSERVAT"Gemini South"Observatory (Gemini-North∣Gemini South)
    OBSID"GS-2014A-SV-412-6"Gemini Observation ID
    CD1_1−3.92777777778E-06partial of first axis coordinate w.r.t. x
    CD1_29.84069992011E-14partial of first axis coordinate w.r.t. y
    CD2_19.8406999118E-14partial of second axis coordinate w.r.t. x
    CD2_23.92777777778E-06partial of second axis coordinate w.r.t. y
    CDELT10.0014Coordinate increment
    CDELT20.0014Coordinate increment
    CRPIX1141.0 x-coordinate of ref pixel [note: first pixel is
    CRPIX2141.0 y-coordinate of ref pixel [note: first pixel is
    CRVAL1269.0887R.A. at ref point
    CRVAL2−21.956075decl. at ref point
    CTYPE1"RA–TAN"First axis is R.A.
    CTYPE2"DEC–TAN"Second axis is decl.
    CUNIT1"deg"Units of data
    CUNIT2"deg"Units of data
    RADESYS"FK5"R.A decl. coordinate system reference
    BSCALE1Linear factor in scaling equation
    BZERO0Zero-point in scaling equation
    WCSAXES3Number of axes in WCS system
    CTYPE3"STOKES"Polarization
    CUNIT3"N/A"Polarizations
    CRVAL31I,Q_phi,U_phi,Q,U,LP_I
    CRPIX30Reference pixel location
    CD3_31Stokes axis: images 0 and 1 give orthogonal pol
    FQ0.002997042052516556avg frac of stell/inst pol removed
    FU0.007443801664624007avg frac of stell/inst pol removed
    TARGET"MWC_275"Target Name
    STOKES"I,Q_phi,U_phi,Q,U,LP_I"data cube Stokes components
    STAR_X140.5star position axis1 (1-based coordinates)
    STAR_Y140.5star position axis2 (1-based coordinates)
    FILTER"J-band "filter band of observation
    ZEROPT1594Jy
    REFMAG6.195ref. mag of star for flux conversion
    FLUXUNIT"mJy/arcseĉ2"pixel flux units
    CALIBFAC4.00356E-08Conversion factor mJy/arcseĉ2/ADU/sec/coadd
    SCALE14.14Pixel Scale mas/pix
    COMMENT⋯Stokes components only take linear pol. into account
    COMMENT⋯All pol. images are stellar pol. subtracted.
    COMMENT⋯2MASS magnitudes used for flux conversion

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    Footnotes

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    10.3847/1538-3881/ac7be4