Alma Survey Of Circumstellar Disks In The Young Stellar Cluster IC 348
D. Ruíz-Rodríguez, L. A. Cieza, J. P. Williams, S. M. Andrews, D. A. Principe, C. Caceres, H. Canovas, S. Casassus, M. R. Schreiber, J. H. Kastner
MMon. Not. R. Astron. Soc. , 1– ?? (2018) Printed 22 May 2018 (MN L A TEX style file v2.2)
ALMA SURVEY OF CIRCUMSTELLAR DISKS IN THE YOUNGSTELLAR CLUSTER IC 348
D. Ruíz-Rodríguez, , (cid:63) L. A. Cieza, , J. P. Williams, S. M. Andrews, D. A. Principe, C. Caceres, , H. Canovas, S. Casassus, , M. R. Schreiber, , , and J. H. Kastner Chester F. Carlson Center for Imaging Science, School of Physics & Astronomy, and Laboratory for Multiwavelength Astrophysics,Rochester Institute of Technology, 54 Lomb Memorial Drive, Rochester NY 14623 USA. Research School of Astronomy and Astrophysics, Australian National University, Canberra, ACT 2611, Australia Millenium Nucleus “Protoplanetary discs in ALMA Early Science", Chile. Núcleo de Astronomía, Facultad de Ingeniería, Universidad Diego Portales, Av. Ejercito 441, Santiago, Chile. Institute for Astronomy, University of Hawaii at Manoa, Honolulu, HI, 96822, USA. Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA, 02138, USA. Department of Physics and Kavli Institute for Astrophysics and Space Research, Massachusetts Institute of Technology, Cambridge, MA 02139, USA. Departamento de Ciencias Fisicas, Facultad de Ciencias Exactas, Universidad Andres Bello. Av. Fernandez Concha 700, Las Condes, Santiago, Chile. European Space Astronomy Centre (ESA), Camino Bajo del Castillo s / n, 28692 Villanueva de la Cañada, Madrid, Spain. Universidad de Chile, Camino el Observatorio 1515, Santiago, Chile. Instituto de Física y Astronomía, Universidad de Valparaíso, Av. Gran Bretaña 1111, 2360102 Valparaíso, Chile. Núcleo Milenio Formación Planetaria - NPF, Universidad de Valparaiso, Av. Gran Bretaña 1111, Valparaiso, Chile.
ABSTRACT
We present a 1.3 mm continuum survey of the young (2-3 Myr) stellar cluster IC 348,which lies at a distance of 310 pc, and is dominated by low-mass stars (M (cid:63) ∼ (cid:12) ).We observed 136 Class II sources (disks that are optically thick in the infrared) at 0.8 (cid:48)(cid:48) (200au) resolution with a 3 σ sensitivity of ∼ dust ∼ ⊕ ). We detect 40 of the targetsand construct a mm-continuum luminosity function. We compare the disk mass distributionin IC 348 to those of younger and older regions, taking into account the dependence on stellarmass. We find a clear evolution in disk masses from 1 to 5-10 Myr. The disk masses in IC 348are significantly lower than those in Taurus (1-3 Myr) and Lupus (1-3 Myr), similar to those ofChamaleon I, (2-3 Myr) and σ Ori (3-5 Myr) and significantly higher than in Upper Scorpius(5 −
10 Myr). About 20 disks in our sample ( ∼
5% of the cluster members) have estimatedmasses (dust + gas) > Jup and hence might be the precursors of giant planets in the cluster.Some of the most massive disks include transition objects with inner opacity holes based ontheir infrared SEDs. From a stacking analysis of the 96 non-detections, we find that thesedisks have a typical dust mass of just (cid:46) ⊕ , even though the vast majority of their infraredSEDs remain optically thick and show little signs of evolution. Such low-mass disks may bethe precursors of the small rocky planets found by Kepler around M-type stars.
Key words:
Circumstellar Disks, Dust and Gas, Interferometry.
The evolution of protoplanetary disks has been studied for decades,and typical disk lifetimes are well established to be ∼ ∼ (cid:63) E-mail:[email protected] planets to be formed (Sicilia-Aguilar et al. 2006). Determining themain process of disk dispersal is not an easy task since severalphysical mechanisms play a role at di ff erent time scales and radii(Alexander et al. 2014), but studying disk properties as a functionof stellar mass and age can shed light on the frequency and loca-tion of forming planets (Mordasini et al. 2012; Alibert, Mordasini& Benz 2011).One important inference from exoplanet surveys is that planetoccurrence generally decreases with increasing planet size: rockyplanets are much more common than gas giants (Howard et al.2012; Burke et al. 2015; Bonfils et al. 2013). Moreover, the cor- c (cid:13) a r X i v : . [ a s t r o - ph . S R ] M a y D. Ruíz-Rodríguez et al. relation between stellar and planet properties indicates that giantplanet occurrence increases with stellar mass at solar metallicity,with a percentage of 3% around M dwarfs ( ∼ (cid:12) ) increasingto 14% around A stars ( ∼ (cid:12) ) (Johnson et al. 2010). These exo-planet correlations are likely to be connected to disk properties asfunctions of stellar mass.Emission from millimeter-sized dust grains in the disk is gen-erally optically thin in the (sub-)millimeter regime; therefore, (sub-)millimeter continuum surveys of disks in star-forming regionswith di ff erent ages ( ∼ −
10 Myr) can trace the distribution of diskmasses as a function of age and stellar mass. This allows us toinvestigate how disk properties and evolution connect to the pop-ulation of planets observed in the field. To exploit this observa-tional potential, Andrews et al. (2013) performed a millimeter con-tinuum survey with the Submillimeter Array (SMA) of the TaurusClass II members (optically thick disks) with spectral types earlierthan M8.5. As a main result, they showed a correlation between themm luminosity (L mm ) and the mass of the host stellar object of theform L mm ∝ M ∗ . − . , which in turn suggests a linear relationshipbetween the masses of the disk and that of the parent star: M dust ∝ M ∗ . Various observational studies of higher sensitivity and res-olution with the Atacama Large Millimeter / submillimeter Array(ALMA) add additional samples in Lupus (1 − − σ Ori (3 − −
10 Myr; Pecaut, Mamajek & Bubar 2012;Barenfeld et al. 2016). A Bayesian linear regression has been thestandard method used to characterize the M dust - M ∗ relations ofthese star-forming regions. Although initially M dust and M ∗ werethought to be linearly correlated in 1 − dust - M ∗ relationship because the limited sensi-tivity implies lower detection rates for late-type stars and browndwarfs, allowing for the possibility of a steeper relation. Indeed,Pascucci et al. (2016) reanalyzed all the submillimeter fluxes andstellar properties available for Taurus, Lupus, and Upper Sco, andfound steeper correlations than linear for these clusters. They alsoobtained a steep dust mass-stellar mass scaling relation in the ∼ dust - M ∗ relation is shared by star-forming regions that are 1-3Myr old (Pascucci et al. 2016). More recently, a similar steepeningof the M dust - M ∗ relation was found by Ansdell et al. (2017) for the σ Ori star-forming region. This steeper relation possibly indicates1) undetected large pebbles or 2) an e ffi cient inward drift in disksaround the lowest-mass stars. In addition, this steepening of theM dust - M ∗ correlation with age suggests a faster decline of circum-stellar dust mass with time in late-type stars. From these relations,at an age of < ∼
10 Myrs, disks around 0.1 and 0.5 M (cid:12) stars mighthave dispersed millimeter-sized grains by factors of 5 and 2.5, re-spectively, faster than earlier-type objects (Pascucci et al. 2016).Following these studies, the IC 348 star-forming region, witha disk fraction of 36% in the IR regime, is an excellent benchmarkto characterize the relationship between the masses of the disk andthat of the host star by comparing to other star-forming regions.In fact, the first millimeter observations of protoplanetary disks inIC348 star-forming region were made by Lee, Williams & Cieza(2011), and with a detection rate of only ∼ Figure 1.
Distribution of stellar spectral types for our sample in the IC 348star-forming region. These targets were selected from Muench et al. (2007)and Lada et al. (2006) and are listed in Table 1 beyond millimeter sizes, resulting in low luminosities at millimeterwavelengths.In this work, we present a 1.3 mm /
230 GHz study of ∼ IC 348 is a rich and compact (2 × ∼
480 members have been identified initially by H α emission (Herbig 1954) and subsequently by optical and IR pho-tometry and spectroscopy (Lada & Lada 1995; Herbig 1998; Luh-man et al. 1998; Luhman 1999; Luhman et al. 2003; Luhman 2003,1999; Luhman, McLeod & Goldenson 2005; Luhman, Esplin &Loutrel 2016). Most of the known T Tauri stars in the IC348 star-forming region have been well studied and spectrally classified(Luhman et al. 2003; Muench et al. 2007): see Figure 1.Our sample was selected specifically from the work of Ladaet al. (2006), whose sample was based on Luhman et al. (2003),and from Muench et al. (2007), a census of 192 candidate YSOs inthe IC 348 nebula, covering a 26.8 (cid:48) × (cid:48) region and centeredat R.A. 03 h m s , Dec. + o (cid:48) (cid:48)(cid:48) . These programs used Spitzer -IRAC photometry to investigate both the frequency and na-ture of the circumstellar disk population in the IC348 cluster on thebasis of the IR SED slope between 3.6 and 8.0 µ m, α . − . µ m . Ingeneral, Lada et al. (2006) and Muench et al. (2007) used α . − . µ m to classify the objects as follows:(i) Class I (protostars): α . − . µ m > − − > α . − . µ m > − / III (anemic disks): − > α . − . µ m > − α . − . µ m < − c (cid:13) , 1– ?? LMA SURVEY OF CIRCUMSTELLAR DISKS IN IC 348 Figure 2.
IR map of IC 348 star-forming region with our ALMA targets.ALLWISE 3-color image with RGB mapped to 22 (W4), 4.6 (W2), and3.4 (W1) µ m. Yellow and blue circles correspond to the sampled selectedfrom Muench et al. (2007) and Lada et al. (2006), respectively. Red crossesindicate the positions of IC 348 members used to estimate a mean clusterdistance of 310 pc ±
20 pc, based on the
Gaia
DR2 parallax data.
Lada et al. (2006) classified as “anemic disks” those objectswith − > α . − . µ m > − Spitzer sources with α . − . µ m values between − − α . − . µ m values, but their 24 µ mfluxes indicate that they are transitional objects with optically thickouter disks (Lada et al. 2006; Espaillat et al. 2012).We note that the standard YSO Class system (Greene et al.1994) is based on the SED slope between ∼ ∼ µ m, butmost IC 348 members lack Spitzer µ m detections. With thesecaveats, our final target list (Table 1) is composed of 136 ClassII disk objects with stellar spectral types in the range of G1 − M9.Figure 2 shows the positions of our targets. Among the objects se-lected, Cl* IC 348 LRL 237, V* V716 Per, Cl* IC 348 LRL 135and Cl* IC 348 LRL 97 are classified by Espaillat et al. (2012) astransitional disks. Our sample also includes Cl* IC 348 LRL 31 andCl* IC 348 LRL 135, which have known close stellar companionsat separations of 38.1 ± ± ALMA observations toward our IC 348 targets were carried outin Band 6 (211-275 GHz) under the project code: 2015.1.01037.S.Our science goal was executed in Cycle 3 with the C40-4 array con-figuration and was observed between the 23rd and 27th June, 2015.The Band 6 continuum observations were conducted with a totalon-source integration time of ∼ ff ective bandwidths of 1.875 GHz centered at 218.0 and233.0 GHz, for a mean frequency of 225.676 GHz ( ∼ σ ) noise level reached is ∼ / beam. We also tar-geted molecular emission lines of CO, CO, and C O (J = − − ) and abandwidth of 117.2 MHz. The ALMA data were reduced using theCommon Astronomy Software Application (CASA) package, ver-sion 4.5.3 (McMullin et al. 2007). Initial calibration (i.e. water va-por radiometer corrections, phase and amplitude calibrations) wasperformed by the ALMA science operations team during qualityassurance. The flux calibrator was J0237 + + + + ∼ +
2, which is close toa natural weighting. Using the uvcontsub routine, we subtracted thecontinuum emission from the spectral windows to extract the CO, CO, and C O spectral line data from the calibrated visibilities.
We searched for 1.3 mm continuum emission centered on the2MASS positions of the 136 targets, listed in Table 1. From the con-tinuum images, we determined the peak flux and rms using the task imstat and thus estimated the signal to noise (S / N, ratio betweenpeak and rms) for each image. Peak fluxes were derived from a 4 (cid:48)(cid:48) radius circle, and the rms from a 4-7 (cid:48)(cid:48) radius annulus centered onthe expected source position. A source with S / N < uvmodelfit routine in CASA and by fitting a point sourcein the uv plane. If the flux density is less than 4 σ , the point sourcefit is applied to the visibilities with the pointing center as a freeparameter. If the flux density is less than 3 σ it is fit with a pointsource with the o ff set position fixed. Table 2 lists integrated fluxdensity (F . ) and rms for non-detected sources.For detections (S / N > uvmodelfit in CASA. This model is centered at the nominal sourceposition and provides the parameters F . , the FWHM along themajor axis, aspect ratio, position angle of the major axis (P.A.), andcoordinate o ff sets ( ∆ α , and ∆ δ ). These parameters are listed in Ta-ble 3. A disadvantage of fitting the brightness profile of a source inthe UV-plane directly is the possibility of including emission froma second source in the fitting process. To avoid any contaminationin the measured flux of each field, we visually inspected the imageplane for pixels with significant brightness ( > σ ). Applying thesemethods, we detect 40 out of the 136 IC 348 targets at > σ sig-nificance. Images of the 40 sources are displayed in Figure 3. Wefind that 10 of the targets are partially resolved, giving P.A. valueswith large uncertainties, and therefore, we do not report those val-ues here. For these objects, the source sizes (deconvolved from thebeam) are listed in Table 3.Using standard approaches (e.g. Hildebrand 1983), the mil-limeter flux can be translated into a disk mass according M dust = F ν d κ ν B ν ( T dust ) , (1) c (cid:13) , 1–, 1–
We searched for 1.3 mm continuum emission centered on the2MASS positions of the 136 targets, listed in Table 1. From the con-tinuum images, we determined the peak flux and rms using the task imstat and thus estimated the signal to noise (S / N, ratio betweenpeak and rms) for each image. Peak fluxes were derived from a 4 (cid:48)(cid:48) radius circle, and the rms from a 4-7 (cid:48)(cid:48) radius annulus centered onthe expected source position. A source with S / N < uvmodelfit routine in CASA and by fitting a point sourcein the uv plane. If the flux density is less than 4 σ , the point sourcefit is applied to the visibilities with the pointing center as a freeparameter. If the flux density is less than 3 σ it is fit with a pointsource with the o ff set position fixed. Table 2 lists integrated fluxdensity (F . ) and rms for non-detected sources.For detections (S / N > uvmodelfit in CASA. This model is centered at the nominal sourceposition and provides the parameters F . , the FWHM along themajor axis, aspect ratio, position angle of the major axis (P.A.), andcoordinate o ff sets ( ∆ α , and ∆ δ ). These parameters are listed in Ta-ble 3. A disadvantage of fitting the brightness profile of a source inthe UV-plane directly is the possibility of including emission froma second source in the fitting process. To avoid any contaminationin the measured flux of each field, we visually inspected the imageplane for pixels with significant brightness ( > σ ). Applying thesemethods, we detect 40 out of the 136 IC 348 targets at > σ sig-nificance. Images of the 40 sources are displayed in Figure 3. Wefind that 10 of the targets are partially resolved, giving P.A. valueswith large uncertainties, and therefore, we do not report those val-ues here. For these objects, the source sizes (deconvolved from thebeam) are listed in Table 3.Using standard approaches (e.g. Hildebrand 1983), the mil-limeter flux can be translated into a disk mass according M dust = F ν d κ ν B ν ( T dust ) , (1) c (cid:13) , 1–, 1– ?? D. Ruíz-Rodríguez et al.
Table 1: Targeted Class II Objects in IC 348.
Source ID Target R.A. Dec. Spec. type Ref. +
32 10 04.88 A2 12 V* V909 Per 03 44 26.03 +
32 04 30.41 G8 13 Cl* IC 348 LRL 13 03 43 59.65 +
32 01 53.98 M0.5 14 V* V926 Per 03 44 44.72 +
32 04 02.48 M0.5 15 Cl* IC 348 LRL 19 03 44 30.82 +
32 09 55.80 A2 16 Cl* IC 348 LRL 26 03 43 56.03 +
32 02 13.21 K7 17 V* V920 Per 03 44 37.88 +
32 08 04.18 K7 18 V* V715 Per 03 44 38.46 +
32 07 35.70 K6 19 V* V712 Per 03 44 37.98 +
32 03 29.66 K6 110 V* V910 Per 03 44 29.73 +
32 10 39.84 K8 111 V* V697 Per 03 44 21.61 +
32 10 37.68 K7 112 Cl* IC 348 LRL 46 03 44 11.62 +
32 03 13.18 – 113 IRAS 03410 + +
32 01 35.50 – 114 Cl* IC 348 LRL 55 03 44 31.37 +
32 00 14.05 M0.5 115* V* V716 Per 03 44 38.54 +
32 08 00.65 M1.25 1, 316 V* V698 Per 03 44 22.29 +
32 05 42.79 K8 117 Cl* IC 348 LRL 63 03 43 58.91 +
32 11 27.07 M1.75 118 Cl* IC 348 LRL 68 03 44 28.51 +
31 59 54.00 M3.5 119 V* V719 Per 03 44 43.77 +
32 10 30.41 M1.25 120 Cl* IC 348 LRL 76 03 44 39.80 +
32 18 04.19 M3.75 121 V* V710 Per 03 44 37.41 +
32 09 00.91 M1 122 V* V922 Per 03 44 39.20 +
32 09 44.90 M2 123* Cl* IC 348 LRL 97 03 44 25.55 +
32 06 17.13 M2.25 1,324 V* V695 Per 03 44 19.24 +
32 07 34.74 M3.75 125 V* V905 Per 03 44 22.32 +
32 12 00.70 M1 126 V* V925 Per 03 44 44.59 +
32 08 12.54 M2 127 V* V919 Per 03 44 37.39 +
32 12 24.20 M2 128 Cl* IC 348 LRL 128 03 44 20.18 +
32 08 56.59 M2 129 Cl* IC 348 LRL 129 03 44 21.30 +
32 11 56.34 M2 130* Cl* IC 348 LRL 135 03 44 39.18 +
32 20 08.93 M4.5 1,331 V* V907 Per 03 44 25.30 +
32 10 12.80 M4.75 132 Cl* IC 348 LRL 140 03 44 35.69 +
32 03 03.54 M3.25 133 Cl* IC 348 LRL 149 03 44 36.98 +
32 08 34.20 M4.75 134 Cl* IC 348 LRL 153 03 44 42.76 +
32 08 33.77 M4.75 135 Cl* IC 348 LRL 156 03 44 06.78 +
32 07 54.09 M4.25 136 V* V902 Per 03 44 18.58 +
32 12 53.08 M2.75 137 Cl* IC 348 LRL 165 03 44 35.46 +
32 08 56.35 M5.25 138 Cl* IC 348 LRL 166A 03 44 42.58 +
32 10 02.50 M4.25 139 Cl* IC 348 LRL 168 03 44 31.35 +
32 10 46.98 M4.25 140 Cl* IC 348 LRL 173 03 44 10.13 +
32 04 04.50 M5.75 141 Cl* IC 348 LRL 192 03 44 23.65 +
32 01 52.69 M4.5 142 V* V713 Per 03 44 38.01 +
32 11 37.03 M4 143 Cl* IC 348 LRL 202 03 44 34.28 +
32 12 40.73 M3.5 144 Cl* IC 348 LRL 203 03 44 18.10 +
32 10 53.44 M0.75 145 Cl* IC 348 LRL 205 03 44 29.80 +
32 00 54.58 M6 146 Cl* IC 348 LRL 214 03 44 07.51 +
32 04 08.81 M4.75 147 Cl* IC 348 LRL 221 03 44 40.24 +
32 09 33.13 M4.5 148 SSTc2d J034431.2 + +
32 05 58.90 M0.5 149* Cl* IC 348 LRL 229 03 44 57.86 +
32 04 01.60 M5.25 1, 350 Cl* IC 348 LRL 237 03 44 23.57 +
32 09 33.88 M5 151 Cl* IC 348 LRL 241 03 44 59.83 +
32 13 31.90 M4.5 152 Cl* IC 348 LRL 248 03 44 35.95 +
32 09 24.31 M5.25 153 Cl* IC 348 LRL 256 03 43 55.27 +
32 07 53.31 M5.75 154 Cl* IC 348 LRL 272 03 44 34.13 +
32 16 35.77 M4.25 155 Cl* IC 348 LRL 276 03 44 09.21 +
32 02 37.68 M0 156 Cl* IC 348 H 149 03 44 34.05 +
32 06 57.05 M7.25 157 Cl* IC 348 LRL 292 03 43 59.87 +
32 04 41.44 M5.75 158 Cl* IC 348 LRL 297 03 44 33.21 +
32 12 57.46 M4.5 159 Cl* IC 348 LRL 300 03 44 38.97 +
32 03 19.69 M5 160 Cl* IC 348 LRL 319 03 45 01.00 +
32 12 22.21 M5.5 161 Cl* IC 348 LRL 324 03 44 45.22 +
32 10 55.75 M5.75 162 Cl* IC 348 LRL 325 03 44 30.06 +
32 08 48.90 M6 163 Cl* IC 348 LRL 334 03 44 26.66 +
32 02 36.32 M5.75 164 Cl* IC 348 LRL 336 03 44 32.37 +
32 03 27.48 M5.5 165 Cl* IC 348 LRL 341 03 44 12.98 +
32 13 15.61 M5.25 166 Cl* IC 348 LRL 366 03 44 35.02 +
32 08 57.34 M4.75 167 Cl* IC 348 LRL 382 03 44 30.96 +
32 02 44.18 M5.5 168 Cl* IC 348 LRL 407 03 45 04.14 +
32 05 04.38 M7 169 Cl* IC 348 LRL 415 03 44 29.97 +
32 09 39.45 M6.5 170 2MASS J03444593 + +
32 03 56.78 M5.75 1 c (cid:13) , 1– ?? LMA SURVEY OF CIRCUMSTELLAR DISKS IN IC 348 Table 1 – Continued
Source Target R.A. Dec. Spect. type Ref.
71 Cl* IC 348 LRL 462 03 44 24.46 +
32 01 43.71 M3 172 Cl* IC 348 LRL 468 03 44 11.07 +
32 01 43.60 M8.25 173 Cl* IC 348 LRL 555 03 44 41.22 +
32 06 27.14 M5.75 174 Cl* IC 348 LRL 603 03 44 33.42 +
32 10 31.50 M8.5 175 [PSZ2003] J034437.6 + +
32 08 32.90 M5.5 176 [PSZ2003] J034426.4 + +
32 08 09.94 M9 177 Cl* IC 348 LRL 690 03 44 36.38 +
32 03 05.40 M8.75 178 Cl* IC 348 LRL 703 03 44 36.62 +
32 03 44.20 M8 179 [PSZ2003] J034433.7 + +
32 05 20.67 M6 180 [PSZ2003] J034433.7 + +
32 05 46.71 M8.75 181 Cl* IC 348 LRL 746 03 44 49.96 +
32 06 14.61 M5 182 [PSZ2003] J034419.7 + +
32 06 45.93 M7 183 Cl* IC 348 LRL 2096 03 44 12.94 +
32 13 24.06 M6 184 [PSZ2003] J034416.2 + +
32 05 40.96 M9 185 Cl* IC 348 TJ 72 03 44 31.98 +
32 11 43.95 G0 186 [BNM2013] 32.03 53 03 44 42.01 +
32 08 59.98 M4.25 187 Cl* IC 348 LRL 8078 03 44 26.68 +
32 08 20.35 M0.5 188 Cl* IC 348 LRL 9024 03 44 35.37 +
32 07 36.24 M0 189 Cl* IC 348 H 110 03 44 25.58 +
32 11 30.24 M2 190 2MASS J03452514 + +
32 09 30.18 M3.75 191 2MASS J03452046 + +
32 06 34.48 M1 192* Cl* IC 348 LRL 31 03 44 18.17 +
32 04 57.04 G1 193* Cl* IC 348 LRL 329 03 44 15.58 +
32 09 21.83 M7.5 194* Cl* IC 348 LRL 67 03 43 44.61 +
32 08 17.76 M0.75 195 2MASS J03435856 + +
32 17 27.53 M3.5(IR) 296 Cl* IC 348 LRL 117 03 43 59.08 +
32 14 21.31 M3.5(IR) 297 2MASS J03442724 + +
32 14 20.98 M3.5(IR) 298 2MASS J03434881 + +
32 15 51.55 M4.5(IR) 299 Cl* IC 348 LRL 179 03 44 34.99 +
32 15 31.15 M3.5(IR) 2100 Cl* IC 348 LRL 199 03 43 57.22 +
32 01 33.90 M6.75(IR) 2101 Cl* IC 348 LRL 215 03 44 28.95 +
32 01 37.85 M3.25(IR) 2102 2MASS J03443112 + +
32 18 48.49 M3.25(IR) 2103* 2MASS J03443468 + +
32 16 00.09 M3.5(IR) 2, 3104 2MASS J03441522 + +
32 19 42.18 M4.75(IR) 2105 2MASS J03442294 + +
32 14 40.43 M5.5(IR) 2106 2MASS J03440599 + +
32 15 32.15 M6.5(IR) 2107 Cl* IC 348 LRL 364 03 44 43.01 +
32 15 59.67 M4.75(IR) 2108 Cl* IC 348 LRL 368 03 44 25.70 +
32 15 49.27 M5.5(IR) 2109 Cl* IC 348 LRL 406 03 43 46.44 +
32 11 05.94 M5.75(IR) 2110 2MASS J03445853 + +
31 58 27.03 M6.5(IR) 2111 2MASS J03432845 + +
32 05 05.82 M4(IR) 2112 Cl* IC 348 LRL 753 03 44 57.62 +
32 06 31.25 XXX 2113 2MASS J03445688 + +
32 20 35.52 M6(IR) 2114 Cl* IC 348 LRL 1379 03 44 52.00 +
31 59 21.50 M9.75 2115 2MASS J03445205 + +
31 58 25.21 M3.5(IR) 2116 Cl* IC 348 LRL 1683 03 44 15.83 +
31 59 36.77 M5.25(IR) 2117 Cl* IC 348 LRL 1707 03 43 47.64 +
32 09 02.56 M7(IR) 2118 2MASS J03451307 + +
32 20 05.32 M5(IR) 2119 2MASS J03442721 + +
32 20 28.82 M5(IR) 2120 2MASS J03435056 + +
32 03 18.00 M8.75(IR) 2121 Cl* IC 348 LRL 1881 03 44 33.79 +
31 58 30.28 M3.75(IR) 2122 2MASS J03432355 + +
32 12 25.82 M4.5(op) 2123 V* V338 Per 03 43 28.20 +
32 01 59.12 M1.75(IR) 2124 Cl* IC 348 LRL 1923 03 44 00.47 +
32 04 32.71 M5(IR) 2125 Cl* IC 348 LRL 1925 03 44 05.77 +
32 00 01.10 M5.5(IR) 2126 EM* LkHA 99 03 45 16.35 +
32 06 19.95 K5(op) 2127 2MASS J03445997 + +
32 22 32.83 M5.25 2128 2MASS J03451782 + +
32 12 05.85 M3.75(op) 2129 2MASS J03431581 + +
32 10 45.53 M4.5(IR) 2130 2MASS J03453563 + +
31 59 54.44 M4.5(IR) 2131 2MASS J03452212 + +
32 05 45.01 M8(IR) 2132 2MASS J03442186 + +
32 17 27.31 M4.75(op) 2133 Cl* IC 348 LRL 22865 03 45 17.65 +
32 07 55.33 L0 2134 Cl* IC 348 LRL 40182 03 45 03.83 +
32 00 23.30 – 2135 Cl* IC 348 LRL 54299 03 43 44.27 +
32 03 42.60 – 2136 2MASS J03451349 + +
32 24 34.71 M4.25 2
Ref.: (1) Lada et al. (2006), (2) Muench et al. (2007), (3) Espaillat et al. (2012). ∗ Transitional Diskc (cid:13) , 1–, 1–
Ref.: (1) Lada et al. (2006), (2) Muench et al. (2007), (3) Espaillat et al. (2012). ∗ Transitional Diskc (cid:13) , 1–, 1– ?? D. Ruíz-Rodríguez et al.
Figure 3. > σ ) in the IC 348 region, see Section 4.1. Each image covers 1.7” X 1.7” size with anaverage beam size of 0.8 (cid:48)(cid:48) X 0.7 (cid:48)(cid:48) . Integrated flux density values are presented at the low-right corner as reported in Table 3.c (cid:13) , 1– ?? LMA SURVEY OF CIRCUMSTELLAR DISKS IN IC 348 Figure 4.
Inferred stellar parameters for IC 348 members (Table 5; Sec.4.2) with theoretical models from Bara ff e et al. (2015) for low mass youngstars overlaid. Solid lines in descending order are 0.5, 1, 2, 3, 5, 10, 20, 50and 100 Myrs isochrones and dashed lines represent the evolutionary tracksin the range of 0.06 and 1.4 M (cid:12) . Blue diamonds represent IC 348 detectedmembers, while red circles correspond to non-detections. where F ν is the integrated flux, d is the distance to the target, B ν ( T dust ) is the Planck function at the average disk temperature, and κ ν is the total opacity. Thus, adopting a distance of 310 pc (Section4.2) and making standard assumptions concerning the disc temper-ature (T dust = κ ν = g − at 1.33 mm;Andrews & Williams (2005, and references therein)), we estimatedisk masses for all detected targets and report them in Table 3.Similarly, the 3 σ upper limits of ∼ Dust ∼ ⊕ . Most of our ALMA targets have fundamental stellar parameterssuch as extinction, stellar masses, luminosity, e ff ective tempera-ture, etc, reported in previous studies. However, not all values havebeen obtained in homogeneous manner, and uncertainties might belarger due to systematic di ff erences in methodology or the adopteddistance to IC 348. Considering that the most recent data releasesof the Gaia
DR2 and Pan-STARRS-1 (PS1) are available, we seekfor uniformity in these estimations. We adopt a uniform distanceof 310 ±
20 pc to all targets based on the
Gaia
DR2 parallax mea-surements (Luri et al. 2018) of 35 targets in our sample that haveDR2 parallax uncertainties smaller than 10%. These objects are in-dicated with red crosses in Figure 2. To estimate the visual extinc-tion (A v ), we use the extinction relations A λ e ff A v listed in Table 4,which are calculated using the extinction law presented in Cardelli,Clayton & Mathis (1989). We use PS1 colours r − z and z − y (Mag-nier et al. 2013), in order of preference, and adopt the relations A z = . r − z ) − ( r − z ) ] and A y = . z − y ) − ( z − y ) ], where ( r − z ) and ( z − y ) are the expected PS1 colours of a main-sequence stargenerated from MARCS synthetic fluxes (Gustafsson et al. 2008).In the special case of Class II objects lacking PS1 photometry inthe necessary bands, we adopted A v from Currie & Kenyon (2009).The stellar luminosities ( L (cid:63) ) of IC 348 members are calcu-lated via the dereddened J -band photometry method of Kenyon &Hartmann (1995) and adopting the distance of 310 pc. We derived Table 2.
Non-Detected Class II Sources in IC 348
Source F . rms Source F . rms[mJy] [mJy beam − ] [mJy] [mJy beam − ] ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± the stellar properties based on the spectral types taken from Luh-man et al. (2003) and Muench et al. (2007) and a conversion fromspectral type to the e ff ective temperature (T e ff ) taken from Pecaut& Mamajek (2013) with uncertainties of ± e ff and L (cid:63) , and assuming all targets are single starsystems, we estimated the stellar masses ( M (cid:63) ) and ages from com-parisons with theoretical pre-MS evolutionary tracks. Masses andages of targets with stellar masses between 0.01 and 1.4 M (cid:12) werederived from models presented in Bara ff e et al. (2015) and stellarmasses > M (cid:12) from the PARSEC evolutionary models (Bressanet al. 2012). Age uncertainties are based mainly on the H-R di-agram placement and the determination of L (cid:63) , incorporating theestimated observational photometry, J-band bolometric correctionand extinction uncertainties. Nevertheless, the dominant sources oferror on the L (cid:63) values are the ∼
10% distance and extinction uncer-tainties (Cieza et al. 2007). Stellar mass uncertainties are dominatedby the ± = e ff , estimated stellar age, A v , L (cid:63) and estimated stellarmass of these objects. Figure 4 shows Bara ff e evolutionary modelswith our IC 348 target selection. The stars are clustered around the2 to 3 Myr isochrones, in agreement with previous age estimatesfor the region. c (cid:13) , 1– ?? D. Ruíz-Rodríguez et al.
Table 3.
Continuum Detections of Class II Sources in IC 348
Source F . (cid:63) rms ∆ α ∆ δ a M Dust [mJy] [mJy beam − ] [ (cid:48)(cid:48) ] [ (cid:48)(cid:48) ] [ (cid:48)(cid:48) ] [M ⊕ ] ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ∗ The elliptical Gaussian model applied to the resolved sources generates five free parameters. Here, wereport: integrated flux density (F . ), FWHM along the major axis (a), right ascension o ff set fromthe phase center ( ∆ α ), and declination o ff set from the phase center ( ∆ δ ).Filter ID A λ e ff A λ e ff A v Zero Point[Å] [Jy]g 4775.6 1.19 3631r 6129.5 0.89 3631i 7484.6 0.67 3631z 8657.8 0.51 3631y 9603.1 0.44 3631
Table 4.
Extinction relations calculated by following the Cardelli, Clayton & Mathis (1989) extinct law with R v = (cid:13) , 1– ?? LMA SURVEY OF CIRCUMSTELLAR DISKS IN IC 348 Table 5: Stellar Properties for Class II Sources in IC 348.
Source ID Log T e ff A v Log L (cid:63) M (cid:63) Source ID Log T e ff A v Log L (cid:63) M (cid:63) [K] [mag] [L (cid:12) ] [M (cid:12) ] [K] [mag] [L (cid:12) ] [M (cid:12) ] + . − . ± + . − . + . − .
69 3.45 + . − . ± + . − . + . − . + . − . ± + . − . + . − .
70 3.45 + . − . ± + . − . + . − . + . − . ± + . − . + . − .
71 3.53 + . − . ± + . − . + . − . + . − . ± + . − . + . − .
72 3.41 + . − . ± + . − . + . − . + . − . ± + . − . + . − .
73 3.45 + . − . ± + . − . + . − . + . − . ± + . − . + . − .
74 3.41 + . − . ± + . − . + . − . + . − . ± + . − . + . − .
75 3.46 + . − . ± + . − . + . − . + . − . ± + . − . + . − .
76 3.39 + . − . ± + . − . + . − . + . − . ± + . − . + . − .
77 3.39 + . − . ± + . − . + . − .
10 3.60 + . − . ± + . − . + . − .
78 3.41 + . − . ± + . − . + . − .
11 3.60 + . − . ± + . − . + . − .
79 3.45 + . − . ± + . − . + . − .
12 3.58 + . − . ± + . − . + . − .
80 3.39 + . − . ± + . − . + . − .
13 3.58 + . − . ± + . − . + . − .
81 3.46 + . − . ± + . − . + . − .
14 3.58 + . − . ± + . − . + . − .
82 3.42 + . − . ± + . − . + . − .
15 3.56 + . − . ± + . − . + . − .
83 3.45 + . − . ± + . − . + . − .
16 3.60 + . − . ± + . − . + . − .
84 3.39 + . − . ± + . − . + . − .
17 3.56 + . − . ± + . − . + . − .
85 3.78 + . − . ± + . − . + . − .
18 3.53 + . − . ± + . − . + . − .
86 3.50 + . − . ± + . − . + . − .
19 3.56 + . − . ± + . − . + . − .
87 3.58 + . − . ± + . − . + . − .
20 3.53 + . − . ± + . − . + . − .
88 3.58 + . − . ± + . − . + . − .
21 3.56 + . − . ± + . − . + . − .
89 3.54 + . − . ± + . − . + . − .
22 3.54 + . − . ± + . − . + . − .
90 3.50 + . − . ± + . − . + . − .
23 3.54 + . − . ± + . − . + . − .
91 3.56 + . − . ± + . − . + . − .
24 3.50 + . − . ± + . − . + . − .
92 3.78 + . − . ± + . − . + . − .
25 3.56 + . − . ± + . − . + . − .
93 3.42 + . − . ± + . − . + . − .
26 3.54 + . − . ± + . − . + . − .
94 3.56 + . − . ± + . − . + . − .
27 3.54 + . − . ± + . − . + . − .
95 3.53 + . − . ± + . − . + . − .
28 3.54 + . − . ± + . − . + . − .
96 3.53 + . − . ± + . − . + . − .
29 3.54 + . − . ± + . − . + . − .
97 3.53 + . − . ± + . − . + . − .
30 3.50 + . − . ± + . − . + . − .
98 3.50 + . − . ± + . − . + . − .
31 3.46 + . − . ± + . − . + . − .
99 3.53 + . − . ± + . − . + . − .
32 3.53 + . − . ± + . − . + . − .
100 3.42 + . − . ± + . − . + . − .
33 3.46 + . − . ± + . − . + . − .
101 3.53 + . − . ± + . − . + . − .
34 3.46 + . − . ± + . − . + . − .
102 3.53 + . − . ± + . − . + . − .
35 3.50 + . − . ± + . − . + . − .
103 3.53 + . − . ± + . − . + . − .
36 3.53 + . − . ± + . − . + . − .
104 3.46 + . − . ± + . − . + . − .
37 3.46 + . − . ± + . − . + . − .
105 3.46 + . − . ± + . − . + . − .
38 3.50 + . − . ± + . − . + . − .
106 3.45 + . − . ± + . − . + . − .
39 3.50 + . − . ± + . − . + . − .
107 3.46 + . − . ± + . − . + . − .
40 3.45 + . − . ± + . − . + . − .
108 3.46 + . − . ± + . − . + . − .
41 3.50 + . − . ± + . − . + . − .
109 3.45 + . − . ± + . − . + . − .
42 3.50 + . − . ± + . − . + . − .
110 3.45 + . − . ± + . − . + . − .
43 3.53 + . − . ± + . − . + . − .
111 3.50 + . − . ± + . − . + . − .
44 3.56 + . − . ± + . − . + . − .
112 3.45 + . − . ± + . − . + . − . c (cid:13) , 1–, 1–
112 3.45 + . − . ± + . − . + . − . c (cid:13) , 1–, 1– ?? D. Ruíz-Rodríguez et al.
Table 5 – Continued
Source ID Log T e ff A v Log L (cid:63) M (cid:63) Source ID Log T e ff A v Log L (cid:63) M (cid:63) [K] [mag] [L (cid:12) ] [M (cid:12) ] [K] [mag] [L (cid:12) ] [M (cid:12) ]
45 3.45 + . − . ± + . − . + . − .
113 3.45 + . − . ± + . − . + . − .
46 3.46 + . − . ± + . − . + . − .
114 3.39 + . − . ± + . − . + . − .
47 3.50 + . − . ± + . − . + . − .
115 3.53 + . − . ± + . − . + . − .
48 3.58 + . − . ± + . − . + . − .
116 3.46 + . − . ± + . − . + . − .
49 3.46 + . − . ± + . − . + . − .
117 3.42 + . − . ± + . − . + . − .
50 3.46 + . − . ± + . − . + . − .
118 3.46 + . − . ± + . − . + . − .
51 3.50 + . − . ± + . − . + . − .
119 3.46 + . − . ± + . − . + . − .
52 3.46 + . − . ± + . − . + . − .
120 3.39 + . − . ± + . − . + . − .
53 3.45 + . − . ± + . − . + . − .
121 3.50 + . − . ± + . − . + . − .
54 3.50 + . − . ± + . − . + . − .
122 3.50 + . − . ± + . − . + . − .
55 3.58 + . − . ± + . − . + . − .
123 3.56 + . − . ± + . − . + . − .
56 3.42 + . − . ± + . − . + . − .
124 3.46 + . − . ± + . − . + . − .
57 3.45 + . − . ± + . − . + . − .
125 3.46 + . − . ± + . − . + . − .
58 3.50 + . − . ± + . − . + . − .
126 3.62 + . − . ± + . − . + . − .
59 3.46 + . − . ± + . − . + . − .
127 3.46 + . − . ± + . − . + . − .
60 3.46 + . − . ± + . − . + . − .
128 3.50 + . − . ± + . − . + . − .
61 3.45 + . − . ± + . − . + . − .
129 3.50 + . − . ± + . − . + . − .
62 3.45 + . − . ± + . − . + . − .
130 3.50 + . − . ± + . − . + . − .
63 3.45 + . − . ± + . − . + . − .
131 3.41 + . − . ± + . − . + . − .
64 3.46 + . − . ± + . − . + . − .
132 3.46 + . − . ± + . − . + . − .
65 3.46 + . − . ± + . − . + . − .
133 - – – –66 3.46 + . − . ± + . − . + . − .
134 3.58 + . − . – – –67 3.46 + . − . ± + . − . + . − .
135 3.58 + . − . – – –68 3.42 + . − . ± + . − . + . − .
136 3.50 + . − . ± + . − . + . − . c (cid:13) , 1– ?? LMA SURVEY OF CIRCUMSTELLAR DISKS IN IC 348 Figure 5.
Spectral energy distributions of the sources detected at 1.3 mm in the IC 348 sample. Red dots show photometric data acquired from the literature;Green dots correspond to our ALMA integrated flux density values; blue lines are the BT-settl spectra model according to the spectral type. A v values usedare in Table 5. The green lines correspond to the median SEDs of K5 − M2 CTTSs calculated by Furlan et al. (2006). The black boxes represent the observedoptical and IR photometry before correcting for extinction. Black triangles are optical photometry upper limitsc (cid:13) , 1– ?? D. Ruíz-Rodríguez et al.
Figure 5.
Continued. c (cid:13) , 1– ?? LMA SURVEY OF CIRCUMSTELLAR DISKS IN IC 348 In Figure 5, we plot the SEDs of all targets detected at 1.3mm, including photometry from PS1 (0.48, 0.62, 0.75, 0.87, 0.96 µ m), 2MASS (1.25, 1.65, 2.22 µ m) and Spitzer / IRAC (3.6, 4.5,5.8, 8 and 24 µ m) (Skrutskie et al. 2006; Evans et al. 2003; Currie& Kenyon 2009). The photometric data were dereddened using theMathis & Cardelli (1990) approach. To calculate the stellar syn-thetic photometry with a fixed temperature T (cid:63) , which is approxi-mated by T e ff , we interpolated the response curves for the set offilters used in the fitting, and used the BT-Settl spectral models forthe corresponding T (cid:63) (Allard 2014). Then, we convolved the filterresponse curves with the synthetic spectra, to match the spectralresolution. Because the PS1, 2MASS, IRAC, and 24 µ m data havephotometric uncertainties between a few percent and 0.1 mag forthe objects investigated here, systematic e ff ects can contribute up to0.1 mag. To account for flux variability of the objects, we added anobservational error of 15%. A multiplicative dilution factor, (cid:16) R (cid:63) d (cid:17) ,relating the central star radius ( R (cid:63) ) and the distance to the object( d ) is used to normalize the optical bands.In Figure 6, we plot millimeter flux as a function of stellarmass. The 8 transition disks in our sample are indicated as redsymbols. Some of these objects are among the most massive disksin the cluster, with disk masses of several M Jup , assuming a stan-dard gas to dust mass ratio of 100. In particular, 3 of the 6 bright-est disks in the entire sample are transition objects based on theirSEDs (Cl* IC 348 LRL 31, Cl* IC 348 LRL 67, and 2MASSJ03443468 + µ m.Three transition disks (Cl* IC 348 LRL 97*, Cl* IC 348 LRL229*, Cl* IC 348 LRL 329) remained undetected. These results fitwell in the scenario proposed by Owen & Clarke (2012) and Ciezaet al. (2012), in which there are at least two types of transition diskswith inner opacity cavities that are the result of distinct processes:1) gas-accreting transition disks that are massive and have large in-ner holes caused by the formation of giant planets, (multiple) lowermass planets or subsequent migration (van der Marel et al. 2018),and 2) non-accreting transition objects with low disk masses thathave inner holes carved by photoevaporation during the final stagesof disk dissipation. The ensemble of undetected sources can be used to estimate thetypical disk mass of the faint sources in IC 348. Initially, westacked di ff erent spectral type sub-groups of these non-detectionsto constrain their average properties. We do not obtain significantdetections in stacked images. Therefore, we stacked the 96 non-detections, after centering each field on the expected stellar posi-tion, to create an average image that has noise which is a factorof ∼ ± σ noise of ourobservations. The 0.14 mJy flux measurement resulting from thestacking exercise suggests that the average dust mass of the disksthat were not individually detected is only ∼ ⊕ . This impliesthat, for most disks in the IC 348 cluster, the amount of millimeter-sized dust that is still available for planet formation is of the orderof the mass of the planet Mars. Kepler has recently found that M- type stars host an average of 2.2 ± ∼ ⊕ and orbital periods of 1.5 to 180 days (Gaidos et al. 2016);therefore, it is expected that most stars in IC 348 should formmultiple rocky planets even though most of the cluster membershave already lost their disks (within the stringent limits imposedby the infrared observations) or have very little dust left. Thus, weconclude that most disks around IC 348 members contain severalEarth masses worth of solids in bodies that are at least several cmin size i. e., large enough to become undetectable by ALMA obser-vations. More significantly, this suggests that these protoplanetarydisks are likely sites of recently formed planetary systems like ourown. In addition, IR emission from disks not detected at mm wave-lengths connotes the existence of small, optically thick disks withextensions of < < ⊕ , which still leaves significant room to ex-plain the observed IR excesses. In fact, a small amount of warmgrains of micron sizes ( < ffi cient to produce theobserved excesses at 10 µ m (e.g. Nagel et al. 2010).In addition, our survey RMS of ∼ / Nof >
4) at 1.3 mm would require 10 × our exposure time. However,we note that objects with a disk mass of ∼ ⊕ are individuallydetected with a S / N of ∼
10 in the Lupus survey thanks to the muchsmaller (150 pc) distance of some of the Lupus PMS stars and, toa lesser extent, the use of a shorter observing wavelength (Ansdellet al. 2016). This implies that at a distance of 150 pc and a sensi-tivity of 0.45 mJy in Band-6, it should possible to detect disks withdust masses of only ∼ ⊕ . Disk properties determine possible planet formation scenarios. In-vestigating basic disk parameters such as mass and size at di ff erentevolutionary stages is thus vital for planet-formation theory. Nearbystar-forming regions like Taurus (1-3 Myr), Lupus (1-3 Myr), Cha I(2-3 Myr), σ Ori (3-5 Myr), and Upper Sco (5-10 Myr) are ideal tar-gets to track evolutionary patterns because the ages of these popu-lations cover the disk dispersal timescale. Recently, (sub-)mm con-tinuum flux surveys of these star-forming regions have shown thatdisk masses decline with age and that there is a strong dependenceof mm-wavelength luminosity on stellar mass (Andrews et al. 2013;Ansdell et al. 2016, 2017; Barenfeld et al. 2016; Pascucci et al.2016). Therefore, in order to compare IC 348 to other regions andinvestigate the evolution of disk masses as a function of stellar age,we need to take into account that disk masses and millimeter de-tection rates depend on spectral types and stellar mass. Figure 8displays the distribution of stellar spectral types for the detectedand non-detected sources, showing the low detection rate at laterspectral types (Tables 2 and 3).Because estimates of stellar masses depend sensitively on in-puts such as distances and theoretical models, here we used the sta-tistical methodology presented in Andrews et al. (2013) based onspectral types. While solar-mass stars evolve in spectral types dur-ing pre-main-sequence stages, lower-mass stars (0.1-0.7 M (cid:12) evolveat almost constant temperature for the first ∼
10 Myr (see evolu-tionary models in Figure 4). This supports the use of spectral typesas a proxy for stellar mass in the mass range of the stars in ourIC 348 sample. Hence, to statistically compare samples from dif-ferent regions, we perform Monte Carlo simulations, whose “ref- c (cid:13) , 1–, 1–
10 Myr (see evolu-tionary models in Figure 4). This supports the use of spectral typesas a proxy for stellar mass in the mass range of the stars in ourIC 348 sample. Hence, to statistically compare samples from dif-ferent regions, we perform Monte Carlo simulations, whose “ref- c (cid:13) , 1–, 1– ?? D. Ruíz-Rodríguez et al.
Figure 6. ∼ Jup assuming a gas to dust mass ratio of 100. -2-1012 D e c . o ff s e t [ " ] -0.05 0.00 0.05 0.10mJy/beam Figure 7.
Stacked image for the 96 non-detections, clearly showing a de-tection at the 6 σ level. erence” sample is IC 348, while a “comparison" sample can beTaurus, Chamaeleon I, Lupus, Upper Sco, or σ Ori. The “compari-son" sample is appropriately scaled to the IC 348 distance (310 pc)and modified for the respective observing wavelengths using the
Figure 8.
Distribution of stellar spectral types for the detected and non-detected sources in IC 348 targeted by our ALMA survey (Tables 2 and3). mean (sub-)millimeter flux ratios observed in Taurus (F λ = F . × (1.3mm / λ ) . ). Upper limit inputs for the “comparison" samplesare as reported in the literature: three times the rms noise of the ob-servations for Taurus, Lupus, Cha I, and Orionis, while the upperlimits in Upper Sco are given by three times the rms noise plus anypositive measured flux density. To construct our simulations, wefirst define a set of spectral type bins ranging from A2 to M6, cor-responding to the distribution of the IC 348 sample, and place thecomparison objects in those bins. Then, disk mm-wave luminosi- c (cid:13) , 1– ?? LMA SURVEY OF CIRCUMSTELLAR DISKS IN IC 348 ties are randomly drawn from the reference region (IC 348) in eachof these spectral type bins, such that the reference and comparisonsamples have the same spectral type distributions. In this manner,we simulate 10 synthetic “reference” disk ensembles that are usedto construct Cumulative Distribution Functions (CDF); see Figure9. Each of these CDFs is compared to the comparison sample toestimate the probability that the two distributions are drawn fromthe same parent population using a censored statistical test (i.e theGehan test: Feigelson & Nelson 1985). The result is a list of 10 such probabilities for each comparison region. The cumulative dis-tributions for these probabilities, f ( p φ ), are also shown in Figure 9(bottom-right panel). The CDFs for the scaled flux densities show that the disks orbit-ing IC 348 stars are fainter on average than disks in Taurus, Lupus,Cham I and σ Ori. Cieza et al. (2015) presented a similar statisticalanalysis based on shallower SCUBA-2 observations of IC 348 andfound that the fluxes in this cluster were slightly lower than in Tau-rus. Here, we confirm that the fluxes in IC 348 are ∼ × fainter thanin Taurus with a very high level of significance: virtually all 10 tests indicate that the probability that disk luminosities in Taurusand IC 348 are drawn from the same parent population is < − .Similarly, we find that the younger Lupus region has a substan-tially brighter distribution compared to the older IC 348 region, atthe (cid:38) σ level. The di ff erence between the luminosity distributionsof IC 348 and Cha I is marginal, with IC 348 being slightly fainterthan Cha I, while, the luminosity distributions of σ Ori and IC 348are statistically indistinguishable ( (cid:46) σ ). In addition, Upper Sco isalso very di ff erent from IC 348 (all tests indicate di ff erences > σ ),but in the opposite sense: the Upper Sco disks are fainter than disksin IC 348, which reflects the fact that the mean dust mass is lowerat the 5-10 Myr age of Upper Sco. In summary, these millime-ter observations trace the population of millimeter / centimeter-sizedgrains at radial distances >
10 au, confirming a significant dispersalprocess in the outer disk over a timescale of ∼ −
10 Myr.Infrared surveys with the
Spitzer Space Telescope , at IRACwavelengths (3.6 − µ m), previously established that the fractionof optically thick dust disk decreases with age, yielding disk frac-tions (%) of 63 ± ± ± ± σ Ori, 36 ± ± <
10 au) from the centralstar. While IR disks observations are very sensitive and typicallyless biased with respect to spectral type, (sub-)millimetre detectionrates are much lower and usually very biased against the lower endof the stellar mass function (M4-M9), making the interpretation ofthe results di ffi cult. Figure 9 (bottom-right panel) compares the disc luminosity distri-butions of the “comparison" and “reference” samples, where p φ isthe probability that the two distributions are drawn from the sameparent population and the vertical green bars indicate the nominal2 σ , 3 σ and 4 σ probabilities. The cumulative distributions derivedfrom the the Peto-Prentice test indicate medians of p φ = − and 5.5 x 10 − for Taurus and Upper Sco, respectively, implying a > σ di ff erence. The Lupus and Cha I samples appear to have adi ff erence of (cid:38) σ in their luminosity distributions, as indicated by medians of p φ = − and p φ = − , respectively.Meanwhile, the σ Ori has a luminosity distribution that is statis-cally indistinguishable ( (cid:46) σ ) from the IC 348 sample, with p φ = − .It is noteworthy that the disc luminosity distribution of our IC348 sample is significantly di ff erent from those of the Taurus andUpper Sco samples. As mentioned above, IC 348 is fainter thanTaurus and Upper Sco is fainter than IC 348, which is not sur-prising, considering their relative ages and their IR disc fractions(Taurus: 63%, Upper Sco: 16%, IC 348: 36%(Ribas et al. 2014)).Also, σ Ori, with an IR disc fraction of 39%, seems to be at anevolutionary stage similar to that of IC 348 in terms of dispersaltimescales.
Recent results for a similar anal-ysis applied to the millimeter surveys of discs towards other star-forming regions (Taurus, Lupus, Cha I, Ori, Upper Sco), also reveala significant decrease of disk masses with “age”, which has been in-terpreted as a signature of evolution. However, ages are di ffi cult todetermine at early times ( <
10 Myr) and are highly dependent onthe adopted distances and theoretical models (Hillenbrand, Bauer-meister & White 2008). Here, we have adopted a “representative"distance value ( ∼ ±
20 pc) based on those IC 348 objects withhigh accuracy
Gaia
DR2 parallaxes. However, our IC 348 sam-ple presents a considerable dispersion in distance, even for objectswith distance uncertainties < ∼ ±
15 pc and the farthest object V* V697 Per islocated at ∼ ±
18 pc. If the distance is less than the adoptedvalue of 310 pc, the target is expected to be even older than 5 Myr.For a distance of ∼
260 pc, the luminosities would decrease by ap-proximately 30% and the inferred age would be around 3-6 Myr(e.g. Ripepi et al. 2014). Similarly, a larger distance would implya younger age for a given target. If the distance is actually ∼ ∼ Gaia distance measurements for all IC 348 members studiedhere, but is its possible that not all targets are actual members ofthe cluster. Revising the membership status of the targets based onthe
Gaia parallaxes and proper motions is beyond the scope of thispaper, but vetting all regions for non-members would certainly beuseful for future ALMA studies of disks in clusters. Here, our Mon-tecarlo simulations are scaled to a distance of 310 pc and we em-phasize that the foregoing comparisons of disc luminosity functionsare highly dependent on the adopted distance and spectral types.In addition, adopting spectral types as a proxy for massalso introduces uncertainties as pre-main-sequence stars, especiallyhigher mass objects, can significantly evolve in spectral type overtime. Given that the mm-emission from disks depends on the hoststellar mass, the results of our statistical analysis can be influencedby the di ff erence in stellar masses at di ff erent evolutionary stagesand spectral types. A commonly used approximation to estimate disk masses is theuse of flux densities in the millimeter wavelength regime, wherethe disk luminosity is proportional to the dust mass (Beckwith et al.1990). In recent years, a Bayesian linear regression approach analy-sis of ALMA surveys of star-forming regions at di ff erent ages haverevealed a positive relationship between dust mass and stellar massbut with a steepening of the M dust - M ∗ relation. (e.g. Andrews et al. c (cid:13) , 1–, 1–
Gaia parallaxes and proper motions is beyond the scope of thispaper, but vetting all regions for non-members would certainly beuseful for future ALMA studies of disks in clusters. Here, our Mon-tecarlo simulations are scaled to a distance of 310 pc and we em-phasize that the foregoing comparisons of disc luminosity functionsare highly dependent on the adopted distance and spectral types.In addition, adopting spectral types as a proxy for massalso introduces uncertainties as pre-main-sequence stars, especiallyhigher mass objects, can significantly evolve in spectral type overtime. Given that the mm-emission from disks depends on the hoststellar mass, the results of our statistical analysis can be influencedby the di ff erence in stellar masses at di ff erent evolutionary stagesand spectral types. A commonly used approximation to estimate disk masses is theuse of flux densities in the millimeter wavelength regime, wherethe disk luminosity is proportional to the dust mass (Beckwith et al.1990). In recent years, a Bayesian linear regression approach analy-sis of ALMA surveys of star-forming regions at di ff erent ages haverevealed a positive relationship between dust mass and stellar massbut with a steepening of the M dust - M ∗ relation. (e.g. Andrews et al. c (cid:13) , 1–, 1– ?? D. Ruíz-Rodríguez et al.
Figure 9.
Cumulative Distribution Functions of the disk luminosities in IC 348 (red) and the 10 synthetic “reference” disk draws, Taurus (1-2 Myr), Lupus(1-3 Myr), Cha I (2-3 Myr), σ Ori (3-5 Myr), and Upper Sco (5-10 Myr), in black colour. At the right bottom, the comparison between the disk luminositydistribution of IC 348 and the “reference” sample shows the probability that IC 348 and the “reference” sample belong to the same population. The verticalgreen bars indicate the nominal 2, 3, and 4 σ probabilities. Figure 10.
Disk dust mass as a function of stellar mass for IC 348 re-gion. Cyan circles represent the detected sources, while yellow trianglesare 4 σ upper limits for non-detections. TDs are displayed by red circlesand triangles, same as Figure 6. The cyan and green solid lines representthe Bayesian linear regression obtained for IC 348 and σ Ori, respectively;see Section 5.3. ∼ σ Ori and UpperSco present a steeper disk mass vs. stellar mass relation (Pascucciet al. 2016; Ansdell et al. 2017). The steepening of this relationwith age has been interpreted in terms of an e ffi cient inward driftof mm-sized grains (Pascucci et al. 2016). However, the param-eters describing the dependence of disk mass on stellar mass arevery sensitive to the mm detection fraction and the treatment of theupper limits. In the case of IC 348, the detection fraction is low( ∼
30 %) and is a strong function of stellar mass. As a result, mostof the detections are restricted to a narrow range of stellar masses.Given these issues, a linear regression fit is not accurate enough toallow a meaningful comparison to other star-forming regions. Nev-ertheless, we report the resulting parameters from the “standard”methodology used in the previous studies mentioned above.Considering all IC 348 sources in our ALMA sample, we de-rive slope and intercept values of β = ± α = ± β and α are the slope and intercept, respectively. Figure10 shows the linear fit obtained from the Bayesian method. Becauseof the di ffi culty in obtaining a reliable fit, we only use σ Ori as acomparison to illustrate di ff erences between our fitting and otherinvestigations with a wider mass range. The linear regression for σ Ori data generated values of 1.95 ± ± β and α , respectively, and an intrinsic scatter value ( δ ) of 0.65 ± c (cid:13) , 1– ?? LMA SURVEY OF CIRCUMSTELLAR DISKS IN IC 348 of δ = ± σ Ori, and Upper Sco (Pascucciet al. 2016; Ansdell et al. 2017). As previously suggested by Pas-cucci et al. (2016), the dispersion can be an intrinsic property ofthe disk population (i.e. disk masses, dust temperatures, and grainsizes) reflected in the diversity of planetary systems.Moreover, M dust measurements are subject to systematic un-certainties in the assumed parameters, such as distances to the star-forming regions. One can hence expect these results for M dust tochange with the availability of
Gaia
DR2 data. In addition, obtain-ing a realistic dust temperature profile is important in the accuracyof estimates of M dust , and to this purpose, it is required a high-resolution data to generate those profiles. However, as shown byTazzari et al. (2017), assuming a constant T dust of ∼
20 K, providesestimates of M dust that are in good agreement with the results ofmore detailed modeling over a wide range of stellar masses, 0.1 to2.0 M (cid:12) . + The brightest millimeter source in our sample, IRAS 03410 + L (cid:12) (Hatchell et al. 2007), was observed previously with the Sub-millimeter Array by Lee, Williams & Cieza (2011). They detecteda bipolar shape in the CO emission, with prominent emissionoutflow lobes and a moderate opening angle. From our observa-tions at a resolution of 0.3 (cid:48)(cid:48) , we are able to estimate position angle(P.A.), mass and kinematics of the of the outflow following the pro-cess presented in Ruíz-Rodríguez et al. (2017). Here, we used the CO emission to correct for the CO optical depth and estimatedthe mass, momentum and kinetic energy of the outflow, see Figure12. Using the C O line, we estimated a systemic velocity of ∼ − . Because CO traces the bipolar and extension cavities ofthe outflow, we drew a line along the rotation axis to estimate aP.A. of ∼ -155 ◦ north through east. Additionally, taking the extentof ∼ (cid:48)(cid:48) ) and maximum speed of the CO emission, weestimated a kinematic age of 1800 yr.To compute the CO mass, we apply the correction factorto all the channels with CO detection above 4 σ . In order to en-sure emission only from the outflow, we built a mask around IRAS03410 + (cid:48)(cid:48) , where emission inside this area wasremoved from the integration. Thus, separating the red- and blue-shifted components, the blue-shifted outflow kinematics were es-timated by integrating channels in the range between 5.0 and 8.0km s − for CO and, 5.0 and 8.0 km s − for CO. The range ofchannels in the redshifted emission is between 9.5 and 17.5 km s − for CO and, 9.5 and 11.5 km s − for CO. To apply the correc-tion factor to all the channels with CO detection, we extrapolatevalues from a parabola fitted to the weighted mean values, wherethe minimum ratio value was fixed at zero velocity. In the fittingprocess, we did not include those data points presented as red dotsin Figure 11, because at these velocities CO starts becoming op-tically thin. The fitted parabola has the form:T T = . + . LSR ) . (2)Table 6 shows the estimates at temperatures of 20 and 50 Kand without correcting for inclination e ff ects. Correcting for the CO optical depth increases the estimated mass of the outflow, themomentum and the kinetic energy by factors of 7.5-43, 5-27, and4-23, respectively, at a temperature of 20 K.
Figure 11.
Intensity ratio between CO and CO function of velocity. Thegreen solid curve is the best-fit second-order polynomial using the blue datapoints, more details in Ruíz-Rodríguez et al. (2017)
IRAS 03410 + α . − . µ m ∼ -0.006 (Lada et al.2006). While the highly dereddened SED peaking in the mid-infrared clearly shows that IRAS 03410 + − M (cid:12) is consistent with the highest mass estimates of previ-ously reported Class 0 and I outflows, after correcting for opticaldepth e ff ects (Dunham et al. 2014). The di ff erences between theseestimates can be attributed to the higher ALMA sensitivities, whichfacilitate the detection of weak and high-velocity emission from theoutflows, thus integrating over high-resolution spectra. In table 6,note that the measured mass of the blue-shifted outflow is a factorof ∼ ff erences in the environment between the cavities. We have observed 136 Class II members of the young stellar clusterIC 348 with ALMA at 1.3 mm. We reach a dust mass sensitivityof 1.3 M ⊕ (3 σ ) and detect a total of 40 disks. The detection rateis a strong function of spectral type, as expected from the knowndependence of disk mass on stellar mass. A stacking analysis ofthe 96 objects that were not individually detected yielded a clear6 σ detection of 0.14 mJy, indicating that these disks have a typicaldust mass of just (cid:46) ⊕ , even though their infrared SEDs remainoptically thick and show little signs of evolution.We compare the disk luminosity function in IC 348 to thosein younger and older regions and see a clear evolution in the dustmasses between 1 and 5-10 Myr. Based on the statistics of extra-solar planets (Gaidos et al. 2016; Howard et al. 2012; Burke et al.2015), a stellar cluster like IC 348 with ∼
400 members dominatedby low-mass stars should form a very small fraction of systems( (cid:46) > Jup in the cluster and the presence of tran-sition disks among this small population. The rest of the membersshould mostly form small rocky planets, consuming most of the pri-mordial dust by the age of the cluster. For the brightest millimetersource in our sample, IRAS 03410 + c (cid:13) , 1–, 1–
400 members dominatedby low-mass stars should form a very small fraction of systems( (cid:46) > Jup in the cluster and the presence of tran-sition disks among this small population. The rest of the membersshould mostly form small rocky planets, consuming most of the pri-mordial dust by the age of the cluster. For the brightest millimetersource in our sample, IRAS 03410 + c (cid:13) , 1–, 1– ?? D. Ruíz-Rodríguez et al.
Table 6.
Mass, Momentum, Luminosity and Kinetic Energy of the Outflow and Envelope
Red shifted (cid:63)
Blue shifted † Isotope Property 50 (K) 20 (K) 50 (K) 20 (K) C O Mass (10 − M (cid:12) ) 23.50 (163.40) 15.89 (110.40) 1.68 (77.73) 1.13 (52.53)Mass loss (10 − M (cid:12) yr − ) 131.18 (911.69) 88.60 (616.14) 9.36 (433.80) 6.33 (293.00)Momentum (10 − M (cid:12) km s − ) 106.00 (556.6) 72.09 (376.21) 1.53 (40.95) 1.04 (27.67)Energy (10 ergs.) 55.38 (220.39) 37.43 (149.00) 0.25 (5.50) 0.16 (3.69)Luminosity (10 − L (cid:12) ) 25.46 (101.33) 17.21 (68.48) 0.11 (2.52) 0.08 (1.70) C O Mass (10 − M (cid:12) ) 0.16 0.11 0.08 0.55Mass loss ( 10 − M (cid:12) yr − ) 0.90 0.60 4.63 3.10Momentum (10 − M (cid:12) km s − ) 0.27 0.18 0.17 0.11Energy (10 ergs.) 0.49 0.33 0.24 0.16Luminosity (10 − L (cid:12) ) 0.02 0.02 0.01 0.01 Blue shifted outflow kinematics were estimated after a cut above 4 σ and integration of channels between 5.0and 8.0 km s − for CO and 5.0 and 8.0 km s − for CO. Red shifted outflow kinematics were estimated with a threshold value above 4 σ and integration of channelsbetween 9.5 and 17.5 km s − for CO, and 9.5 and 11.5 km s − for CO. Parameters inside the parentheses correspond to the computed values after applying the correction factors foroptical depth e ff ects to all the channels with CO detection above 4 σ . Figure 12. CO intensity map (moment-0) of IRAS 03410 + − . Black contours showthe 1.3 mm continuum emission around IRAS 03410 + × rms (0.15 mJy beam − ). Green and Red contours show the blue-and red- shifted moment-0 of the CO line, respectively, at 20, 40, 80, 160x 3 σ levels. These blue- and red- shifted intensity maps are integrated overthe velocity range of 5.0 to 8 km s − and 9.5 to 17.5 km s − , respectively.The synthesized beam of 0.77 arcsec × = ± (cid:48)(cid:48) ) ofthe central object. P.A., mass and kinematics of the outflow. These estimates are char-acteristic of a Class-I type object.
ACKNOWLEDGMENTS
This paper makes use of the following ALMA data:ADS / JAO.ALMA No. 2015.1.01037.S. ALMA is a partner-ship of ESO (representing its member states), NSF (USA) andNINS (Japan), together with NRC (Canada), NSC and ASIAA(Taiwan), and KASI (Republic of Korea), in cooperation with the Republic of Chile. The Joint ALMA Observatory is operated byESO, AUI / NRAO and NAOJ. The National Radio AstronomyObservatory is a facility of the National Science Foundationoperated under cooperative agreement by Associated Universities,Inc. D.R. acknowledges support from NASA Exoplanets programgrant NNX16AB43G. L.A.C., acknowledges support from the Mil-lennium Science Initiative (Chilean Ministry of Economy), throughgrant Nucleus RC130007. L.A.C. was also supported by FONDE-CYT grant number 1171246. D. P. recognizes support by theNational Aeronautics and Space Administration through ChandraAward Number GO6-17013A issued by the Chandra X-ray Obser-vatory Center, which is operated by the Smithsonian AstrophysicalObservatory for and on behalf of the National Aeronautics SpaceAdministration under contract NAS8-03060. C.C. acknowledgessupport from project CONICYT PAI / Concurso Nacional Insercionen la Academia, convocatoria 2015, folio 79150049, and from ICMNucleo Milenio de Formacion Planetaria, NPF. SC acknowledgessupport from FONDECYT grant 1171624. c (cid:13) , 1– ?? LMA SURVEY OF CIRCUMSTELLAR DISKS IN IC 348 REFERENCES
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