The VLA/ALMA Nascent Disk and Multiplicity (VANDAM) Survey of Perseus Protostars. VI. Characterizing the Formation Mechanism for Close Multiple Systems

John J. Tobin; Leslie W. Looney; Zhi-Yun Li; Sarah I. Sadavoy; Michael M. Dunham; Dominique Segura-Cox; Kaitlin Kratter; Claire J. Chandler; Carl Melis; Robert J. Harris; Laura Perez

doi:10.3847/1538-4357/aae1f7

1. Introduction

Star formation typically occurs within dense clouds of molecular gas in the interstellar medium. These clouds of dense gas collapse to form stars when their self-gravity dominates over other sources of support (e.g., thermal pressure, magnetic fields, turbulence; McKee & Ostriker 2007), forming either a single or multiple star system. Nearly half of Sun-like stars (in terms of stellar mass) are found in binary or higher-order multiple systems (Duquennoy & Mayor 1991; Raghavan et al. 2010) with typical separations of ∼50 au. The frequency of stellar multiplicity strongly depends on stellar mass. Stars more massive than the Sun have a higher fraction of multiplicity, and stars less massive than the Sun have a lower degree of multiplicity; see Duchêne & Kraus (2013) for a recent review. Thus, it is clear that multiple star formation is a common outcome of star formation at all masses, and a comprehensive understanding of the star formation process must also account for multiplicity.

There are two favored routes to explain the formation of multiple star systems: disk fragmentation due to gravitational instability (Adams et al. 1989; Bonnell & Bate 1994; Stamatellos & Whitworth 2009; Kratter et al. 2010) and turbulent fragmentation of the molecular cloud (Padoan & Nordlund 2002; Offner et al. 2010). Disk fragmentation will preferentially result in the formation of close (<600 au) multiple star systems and requires the existence of a rotationally supported disk around the primary star. Turbulent fragmentation can result in the formation of both wide and close multiple systems. The initial protostars form with separations ∼1000 au (or larger), and depending on their relative motions and masses they can migrate closer together, remain at wide separations, or drift further apart (Offner et al. 2010; Sadavoy & Stahler 2017). Previously, it was also thought that envelope rotation could play a role in binary/multiple formation (e.g., Burkert & Bodenheimer 1993), but this is no longer favored given the lack of ordered rotation for most star-forming cores (Tobin et al. 2011; Storm et al. 2014; Fernández-López et al. 2014) and the apparently random orientations of angular momenta (traced by outflow directions) for the wide companion protostars (Lee et al. 2016). A key prediction of turbulent fragmentation is that the seeds for multiplicity are produced in the pre-stellar phase. Early Atacama Large Millimeter/submillimeter Array (ALMA) survey results suggest tentative agreement with these predictions (Pineda et al. 2015; Dunham et al. 2016; Kirk et al. 2017), but better statistics are still needed.

Despite the excellent statistics offered by studying field stars, the properties of field star multiples cannot alone reveal the origin of multiplicity. This is because these systems are observed in the present epoch and represent the culmination of Myr to Gyr of dynamical evolution, and to understand the origin of multiplicity, it must be characterized during or shortly after the onset of star formation during the protostellar phase.

Therefore, to gain a clearer picture of multiple star formation, protostars in the earliest phase of the star formation process, known as the Class 0 phase (André et al. 1993), must be observed. During this phase the forming protostar is enshrouded in a dense, infalling envelope of gas and dust. The high levels of obscuration, however, made searching for multiplicity toward the Class 0 protostars difficult in the infrared, especially at scales less than 1000 au. Thus, centimeter and millimeter interferometry were necessary to examine the formation of such systems. Early searches detected a number of wide multiple systems (>600 au) and a few close (<600 au) multiple systems (Rodríguez et al. 1998; Brown et al. 2000; Looney et al. 2000). Later surveys aimed to improve upon the statistics, and Chen et al. (2013) characterized a number of wide systems in the Class 0 phase, finding that most (67%) Class 0 protostars may begin their lives as part of a wide multiple system. Maury et al. (2010) and Tobin et al. (2013, 2015b) aimed to expand the characterization of close multiples, but samples were still too small to provide statistically significant results.

Several efforts were also made to characterize multiplicity in the late protostellar phase (Duchêne et al. 2007; Connelley et al. 2008), the so-called Class I protostars (Dunham et al. 2014). However, even these systems may be too evolved to retain signatures of their formation mechanism given that enough enshrouding material has been accreted/dispersed to enable their detection in the near-infrared.

The VLA/ALMA Nascent Disk and Multiplicity (VANDAM) Survey (Tobin et al. 2015a, 2016b; Segura-Cox et al. 2016) unlocked the distribution of protostellar multiples on scales less than 600 au (down to ∼20 au), thereby sampling the peak of the field solar-type separation distribution. This survey, conducted with the NSF's Karl G. Jansky Very Large Array (VLA) observed all known Class 0 and Class I protostars in the Perseus star-forming region, 37 of which are Class 0 protostars (first hydrostatic cores and very low luminosity objects included), 8 are Class 0/I protostars, and 37 are Class I protostars (flat spectrum included) (Tobin et al. 2016b). In addition, 12 Class II sources (pre-main sequence stars with disks) that were bright in the far-infrared were also observed. We adopt a distance of 300 pc as an average distance for the Perseus molecular cloud (Ortiz-León et al. 2018). There may be a distance gradient across the cloud (Ortiz-León et al. 2018; Zucker et al. 2018) with its eastern part (IC348) at ∼320 pc, central part (NGC 1333) at ∼293 pc, and western part (L1448 and L1451) at ∼280 pc, but 300 pc is adopted for simplicity because these distances are consistent with each other when considering the uncertainties of current Gaia measurements. We note that these revised distances are in disagreement with the maser parallaxes from Hirota et al. (2008, 2011) which indicated ∼230 pc. This 230 pc distance was used in Tobin et al. (2016a, 2016b) to determine physical separations and convert continuum flux density to mass. From this sample, 18 systems were identified as multiples with companion separations less than 600 au, and 16 of these were new detections by the VANDAM survey.

The VANDAM survey greatly increased the detection statistics for protostars on scales less than 600 au, but the formation mechanism(s) of the multiple systems was (were) difficult to quantify because the VLA survey was continuum-only. Tobin et al. (2016b) argued that systems with separations beyond ∼500 au most likely formed from turbulent fragmentation (Offner et al. 2016); this conclusion was further supported by observations of outflow directions by Lee et al. (2016). For the close multiples detected, disk fragmentation was argued to be the most likely formation mechanism for 17 out of the 18¹² based on their proximity, circumbinary material around some sources, and the double-peaked nature of the separation distribution. Nevertheless, it remains a possibility that the close multiple systems initially formed at wide separations and migrated inward (Offner et al. 2010; Sadavoy & Stahler 2017). However, definitive evidence in favor of disk fragmentation would be finding them embedded within a common, circumbinary/multiple structure with Keplerian rotation. To search for the presence or lack of such a signature, we observed the sample of 17 (out of 18) close multiple sources in Perseus with ALMA; an additional close multiple system (IRAS 03282+3035/Per-emb-5) was discovered after the program was accepted. ALMA enables simultaneous observation of both the dust continuum to look for circum-multiple structure and molecular lines to search for rotation signatures in the surrounding gas (e.g., Takakuwa et al. 2014, 2015; Tobin et al. 2016a).

The results presented here reveal the circum-multiple environment around each close system, such as whether or not there is rotation at small -scales and the presence or absence of surrounding material. This paper is structured as follows: we discuss the observations and data reduction in Section 2, the results from dust continuum observations and molecular line mapping in Section 3, the implications of the observations in Section 4, and present our conclusions in Section 5.

2. Observations and Data Reduction

The Perseus multiple systems were observed with ALMA during Cycle 2 on 2015 September 27 with 33 antennas operating and sampling baselines between 32 and 2000 m. The observations were executed within a 1.8 hr block and the total time spent on each source was ∼2.9 min. The precipitable water vapor was ∼0.7 mm throughout the observing session. The phase calibrator was J0319+4130 (3C84), the bandpass calibrator was J0237+2848, and the amplitude and absolute flux calibrator was the monitored quasar J0238+166. The absolute flux calibration accuracy is expected to be better than 10%; however, we do not include this uncertainty in our calculations and utilize statistical uncertainties throughout the paper. The correlator was configured to observe a 2 GHz continuum band centered at 232.5 GHz and observed in TDM mode with 128 channels. The other three basebands were allocated to 60 MHz windows with 1960 channels (0.083 km s⁻¹ velocity resolution) each and centered on the following molecular transitions: ¹²CO ( $J=2\to 1$ ), ¹³CO ( $J=2\to 1$ ), C¹⁸O ( $J=2\to 1$ ), SO ( ${J}_{N}={6}_{5}\to {5}_{4}$ ), and H₂CO ( $J={3}_{03}\to {2}_{02}$ ).

The raw visibility data were manually reduced by the North American ARC staff using CASA version 4.5.0, and we additionally performed self-calibration on the continuum data to increase the signal-to-noise ratio (S/N). When possible, we performed two rounds of phase self-calibration, first with solution intervals that encompassed the length of an entire on-source scan, then the second round utilized either 12.1 or 6.05 s solution intervals. The shortest possible solution interval was 6.05 s, corresponding to the length of a single integration. Following phase self-calibration, we performed amplitude self-calibration using the solution normalization option, which normalizes the amplitude solutions to around a factor 1.0 such that the flux density scaling is not altered. Following the completion of self-calibration on the continuum, the solutions were also applied to the spectral line bands. Self-calibration was possible for 16 out of 17 fields observed; only Per-emb-48 could not be self-calibrated due to insufficient S/N. The resultant noise in the 1.3 mm continuum was ∼0.14 mJy beam⁻¹ and ∼15 mJy beam⁻¹ in the 0.25 km s⁻¹ channels for the spectral line observations. The non-self-calibrated continuum images had a noise level of ∼1 mJy beam⁻¹.

The data were imaged using the clean task within CASA 4.5.0; the ALMA images shown in Figures 1 and 2 for each target were generated using Briggs weighting with a robust parameter of 0.5. Within the clean task, we only include data at uv-distances >50 kλ to mitigate striping in the images from large-scale emission detected on the shortest baseline, which could not be properly imaged. The ¹³CO ( $J=2\to 1$ ) and C¹⁸O ( $J=2\to 1$ ) images were generated with natural weighting, tapering at 500 kλ, and also only using the data having uv-distances >50 kλ. Tapering reduces the weight of longer-baseline data in the deconvolution to facilitate the detection of larger, lower surface brightness structures. Typical beam sizes of the continuum and molecular line images are 0 farcs 27 × 0 farcs 16 (81 au × 48 au) and 0 farcs 35 × 0 farcs 25 (108 au × 76 au), respectively.

**Figure 1.** ALMA images of Class 0 multiple protostar systems in Perseus at 1.3 mm. The white or black crosses mark the VLA source positions in each image. A 1'' scale bar is also drawn in each panel denoting 300 au. The beam of each image is drawn in the lower right corner, corresponding to approximately 027 × 017 (81 au × 51 au). The noise level in each image is approximately 0.14 mJy beam⁻¹, but this varies somewhat between sources depending on dynamic range limits. The approximate outflow directions (when known) are drawn in the lower right corner with the red and blue arrow directions corresponding to the orientation of the outflow. Note that the outflows apparently originate from the bright continuum peaks, but the arrows are drawn offset for clarity.
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farcs — **Figure 1.** ALMA images of Class 0 multiple protostar systems in Perseus at 1.3 mm. The white or black crosses mark the VLA source positions in each image. A 1'' scale bar is also drawn in each panel denoting 300 au. The beam of each image is drawn in the lower right corner, corresponding to approximately 027 × 017 (81 au × 51 au). The noise level in each image is approximately 0.14 mJy beam⁻¹, but this varies somewhat between sources depending on dynamic range limits. The approximate outflow directions (when known) are drawn in the lower right corner with the red and blue arrow directions corresponding to the orientation of the outflow. Note that the outflows apparently originate from the bright continuum peaks, but the arrows are drawn offset for clarity.
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**Figure 2.** Same as Figure 1, except for Class I multiple systems in Perseus.
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We note that the phase was very unstable during the course of our observations, and the phase calibrator (J0319+4130) was at ∼10° lower elevation than our science targets. As such, there are systematic position offsets in each field up to ∼0 farcs 1. Thus, the most precisely measured positions for the observed protostars in Perseus are those listed in Tobin et al. (2016b) from the VLA, despite the higher S/N offered by the ALMA observations.

2.1. Data Analysis

With the observed sample of protostellar multiple systems, one of the key observables is the flux density of the resolved sources and that of the surrounding extended structure. We measured the flux densities toward the binary and wide companion systems that fell within the primary beam using Gaussian fitting. First, we fit Gaussians to all compact, individual components that were resolved and/or marginally resolved in the data. The Gaussians were fit using the imfit task of CASA 4.7.2. The integrated flux densities and uncertainties, source positions, and Gaussian parameters from imfit are given in Table 1. Some sources have clear non-Gaussian, extended components. In these cases, we used casaviewer to draw a polygon region encompassing all the extended emission from the sources (out to ∼2'', where we still recover all flux) and measured the flux density within the polygon; this value is given as extended flux in Table 1. For binary systems, there were many cases where the extended flux was greater than the combined flux densities of the Gaussian fits to the compact sources. This indicates that there is a significant amount of emission in extended, non-Gaussian components.

Table 1. Source List and Gaussian Fitting Results

Source	R.A.	Decl.	Int. Flux	Peak Flux	θ_maj × θ_min	PA	Deconv. θ_maj	Deconv. θ_min	Deconv. PA	Extended Flux	1.3–9.1 mm^a
	(J2000)	(J2000)	(mJy)	(mJy beam⁻¹)	('' × '')	(°)	(arcsec)	(arcsec)	(°)	(mJy)	Spectral Index
Per-emb-12-A	3:29:10.53	31:13:30.99	1503.7 ± 29.2	180.7 ± 3.2	0.77 × 0.64	50.9	0.73 ± 0.02	0.61 ± 0.02	58.2 ± 6.1	3270 ± 140.0	2.9 – 2.7 ± 0.07
Per-emb-12-B	3:29:10.43	31:13:32.15	434.2 ± 7.9	128.6 ± 3.2	0.49 × 0.43	119.9	0.46 ± 0.02	0.33 ± 0.02	116.4 ± 9.8	−99.0 ± −99.0	3.7 – 3.3 ± 0.08
Per-emb-17-A	3:27:39.10	30:13:03.08	26.1 ± 0.5	22.9 ± 0.2	0.28 × 0.17	21.0	0.11 ± 0.01	0.062 ± 0.01	15.0 ± 8.8	55.0 ± 7.0	2.2 – 2.1 ± 0.07
Per-emb-17-B	3:27:39.11	30:13:02.85	20.4 ± 0.5	15.9 ± 0.2	0.31 × 0.2	18.6	0.16 ± 0.01	0.11 ± 0.01	7.1 ± 11.3	−99.0 ± −99.0	2.8 ± 0.09
Per-emb-18	3:29:11.26	31:18:31.03	145.2 ± 1.4	29.0 ± 0.3	0.74 × 0.31	72.3	0.70 ± 0.01	0.19 ± 0.01	75.4 ± 0.3	150.0 ± 9.0	3.0 – 2.8 ± 0.07
Per-emb-21	3:29:10.66	31:18:20.10	49.3 ± 0.5	44.3 ± 0.3	0.29 × 0.18	21.7	0.083 ± 0.01	0.065 ± 0.01	32.4 ± 17.8	53.0 ± 4.0	2.4 ± 0.08
Per-emb-22-A	3:25:22.40	30:45:13.21	33.6 ± 0.7	18.8 ± 0.2	0.36 × 0.23	35.9	0.27 ± 0.01	0.16 ± 0.01	49.7 ± 3.4	74.0 ± 12.0	2.6 – 2.2 ± 0.08
Per-emb-22-B	3:25:22.35	30:45:13.14	10.1 ± 0.2	8.5 ± 0.2	0.26 × 0.15	20.6	0.0 ± 0.0	0.0 ± 0.0	0.0 ± 0.0	−99.0 ± −99.0	2.7 – 2.2 ± 0.09
Per-emb-27-A	3:28:55.57	31:14:36.95	162.5 ± 2.1	108.5 ± 0.8	0.31 × 0.21	24.1	0.17 ± 0.01	0.14 ± 0.01	46.5 ± 10.0	290.0 ± 27.0	2.5 – 2.4 ± 0.07
Per-emb-27-B	3:28:55.56	31:14:36.38	26.8 ± 2.5	15.4 ± 0.8	0.40 × 0.22	25.4	0.29 ± 0.05	0.15 ± 0.02	29.0 ± 10.8	−99.0 ± −99.0	2.4 ± 0.09
Per-emb-2	3:32:17.92	30:49:47.85	571.8 ± 9.7	20.5 ± 0.7	1.22 × 0.83	173.8	1.20 ± 0.02	0.81 ± 0.01	172.6 ± 1.8	570.0 ± 34.0	2.9 ± 0.07
Per-emb-33-A	3:25:36.32	30:45:14.84	6.5 ± 0.8	22.4 ± 0.5	0.25 × 0.15	31.5	0.0 ± 0.0	0.0 ± 0.0	0.0 ± 0.0	740.0 ± 32.0	2.1 – 1.7 ± 0.1
Per-emb-33-B	3:25:36.31	30:45:15.09	31.9 ± 1.4	26.5 ± 0.5	0.39 × 0.21	11.9	0.28 ± 0.02	0.13 ± 0.01	5.7 ± 4.6	−99.0 ± −99.0	2.3 ± 0.08
Per-emb-33-C	3:25:36.38	30:45:14.64	163.0 ± 1.8	79.4 ± 0.5	0.39 × 0.28	21.2	0.27 ± 0.01	0.22 ± 0.01	23.0 ± 2.8	−99.0 ± −99.0	2.9 ± 0.08
L1448NW-A	3:25:35.67	30:45:34.08	49.8 ± 0.5	39.7 ± 0.2	0.31 × 0.19	23.1	0.16 ± 0.01	0.075 ± 0.01	33.6 ± 1.9	82.0 ± 7.0	2.5 – 2.3 ± 0.07
L1448NW-B	3:25:35.68	30:45:34.34	14.8 ± 0.2	19.2 ± 0.2	0.27 × 0.17	18.1	0.0 ± 0.0	0.037 ± 0.0	18.2 ± 0.0	−99.0 ± −99.0	1.9 ± 0.08
Per-emb-44-A	3:29:03.76	31:16:03.81	99.0 ± 1.1	80.5 ± 0.5	0.32 × 0.19	20.7	0.16 ± 0.01	0.093 ± 0.01	23.9 ± 2.3	400.0 ± 19.0	2.4 – 2.2 ± 0.07
Per-emb-44-B	3:29:03.74	31:16:03.80	158.1 ± 1.8	74.9 ± 0.5	0.42 × 0.26	27.3	0.32 ± 0.01	0.20 ± 0.01	33.6 ± 1.6	−99.0 ± −99.0	3.1 – 3.0 ± 0.08
SVS13B	3:29:03.07	31:15:51.72	222.2 ± 2.2	73.5 ± 0.5	0.43 × 0.35	53.0	0.36 ± 0.01	0.26 ± 0.01	75.2 ± 2.5	270.0 ± 28.0	2.7 ± 0.07
SVS13A2	3:29:03.38	31:16:01.62	13.9 ± 1.0	12.8 ± 0.5	0.29 × 0.17	18.1	0.11 ± 0.05	0.044 ± 0.03	10.5 ± 79.8	1.3 ± 4.3	2.2 – 2.0 ± 0.09
RAC1999 VLA20	3:29:04.25	31:16:09.14	4.4 ± 1.0	3.7 ± 0.5	0.28 × 0.19	13.7	0.11 ± 0.1	0.026 ± 0.1	136.5 ± 85.1	4.5 ± 3.0	1.7 ± 0.2
Per-emb-8	3:44:43.98	32:01:35.15	102.8 ± 2.7	64.4 ± 0.3	0.32 × 0.24	24.5	0.17 ± 0.01	0.14 ± 0.01	103.8 ± 4.7	120.0 ± 9.3	2.8 – 2.4 ± 0.07
Per-emb-8 Inner Disk	3:44:43.98	32:01:35.18	89.1 ± 0.7	64.4 ± 0.3	0.30 × 0.23	19.5	0.16 ± 0.01	0.094 ± 0.01	120.9 ± 2.6	120.0 ± 9.3	2.7 – 2.3 ± 0.07
Per-emb-8 Outer Disk	3:44:43.98	32:01:35.15	34.0 ± 2.2	64.4 ± 0.3	1.25 × 0.38	56.5	1.22 ± 0.09	0.32 ± 0.03	57.5 ± 1.6	−99.0 ± −99.0	⋯
Per-emb-35-A	3:28:37.09	31:13:30.74	22.6 ± 0.4	20.5 ± 0.2	0.29 × 0.17	21.6	0.085 ± 0.01	0.054 ± 0.01	36.4 ± 19.6	−99.0 ± −99.0	2.8 – 2.1 ± 0.08
Per-emb-35-B	3:28:37.22	31:13:31.70	16.8 ± 0.4	15.8 ± 0.2	0.28 × 0.16	22.3	0.078 ± 0.02	0.021 ± 0.02	40.3 ± 11.0	−99.0 ± −99.0	2.3 ± 0.08
Per-emb-36-A	3:28:57.37	31:14:15.79	129.2 ± 0.9	88.6 ± 0.4	0.29 × 0.23	22.8	0.16 ± 0.002	0.091 ± 0.01	105.7 ± 2.0	140. ± 5.0	2.4 – 2.3 ± 0.07
Per-emb-36-B	3:28:57.37	31:14:16.11	12.3 ± 0.7	29.6 ± 0.4	0.27 × 0.18	18.9	0.0 ± 0.01	0.0 ± 0.07	0.0 ± 0.0	−99.0 ± −99.0	2.0 ± 0.08
Per-emb-40-A	3:33:16.67	31:07:54.84	15.9 ± 0.3	14.1 ± 0.5	0.30 × 0.17	21.5	0.12 ± 0.01	0.04 ± 0.01	24.1 ± 3.6	17.0 ± 2.0	2.4 – 2.0 ± 0.08
Per-emb-40-B	3:33:16.69	31:07:55.20	0.9 ± 0.1	2.2 ± 0.5	0.28 × 0.16	20.8	0.0 ± 0.0	0.0 ± 0.0	0.0 ± 0.0	−99.0 ± −99.0	1.6 ± 0.15
Per-emb-48-A	3:27:38.28	30:13:58.60	4.0 ± 0.4	2.7 ± 0.1	0.36 × 0.22	21.4	0.23 ± 0.04	0.13 ± 0.02	23.5 ± 18.7	4.5 ± 2.0	2.0 ± 0.14
Per-emb-48-B	⋯	⋯	<0.4	<0.4 ± 0.1	0.36 × 0.22	21.4	0.23 ± 0.04	0.13 ± 0.02	23.5 ± 18.7	4.5 ± 2.0	<1.3
Per-emb-49-A	3:29:12.96	31:18:14.33	18.3 ± 0.4	15.5 ± 0.2	0.28 × 0.18	14.6	0.1 ± 0.01	0.052 ± 0.03	146.2 ± 12.3	20.0 ± 4.3	2.1 ± 0.08
Per-emb-49-B	3:29:12.98	31:18:14.44	4.9 ± 0.3	5.0 ± 0.2	0.26 × 0.16	20.8	0.0 ± 0.0	0.0 ± 0.0	0.0 ± 0.0	−99.0 ± −99.0	2.0 – 1.9 ± 0.1
Per-emb-55-A	3:44:43.29	32:01:31.19	3.3 ± 0.5	3.2 ± 0.3	0.29 × 0.18	25.3	0.078 ± 0.09	0.027 ± 0.06	79.2 ± 70.5	−99.0 ± −99.0	1.8 ± 0.11
Per-emb-55-B	3:44:43.33	32:01:31.65	0.5 ± 0.3	0.5 ± 0.3	0.29 × 0.17	22.6	0.052 ± 0.0	0.04 ± 0.0	17.5 ± 0.0	−99.0 ± −99.0	1.7 – 1.0 ± 0.32
L1448IRS1-A	3:25:09.45	30:46:21.92	73.4 ± 0.9	48.4 ± 0.4	0.33 × 0.2	23.4	0.21 ± 0.01	0.10 ± 0.01	29.3 ± 1.6	78.7 ± 14.0	2.5 – 2.3 ± 0.07
L1448IRS1-B	3:25:09.41	30:46:20.61	5.5 ± 0.7	5.5 ± 0.4	0.25 × 0.17	19.5	0.0 ± 0.0	0.0 ± 0.0	0.0 ± 0.0	−99.0 ± −99.0	2.2 ± 0.13
EDJ2009-269-A	3:30:44.02	30:32:46.82	11.1 ± 0.3	10.6 ± 0.2	0.26 × 0.17	18.5	0.043 ± 0.03	0.039 ± 0.01	147.3 ± 67.5	−99.0 ± −99.0	2.1 – 2.0 ± 0.1
EDJ2009-269-B	3:30:43.98	30:32:46.59	7.9 ± 0.3	7.5 ± 0.2	0.26 × 0.17	19.2	0.0 ± 0.07	0.0 ± 0.05	0.0 ± 0.0	−99.0 ± −99.0	1.9 ± 0.1
L1448IRS3A	3:25:36.50	30:45:21.84	101.0 ± 2.5	29.6 ± 0.5	0.75 × 0.39	136.5	0.73 ± 0.03	0.30 ± 0.02	134.1 ± 2.0	110.0 ± 9.0	2.7 – 2.4 ± 0.07
L1448IRS3A Inner Disk	3:25:36.50	30:45:21.78	165.4 ± 0.8	29.6 ± 0.5	0.26 × 0.16	20.0	0.0 ± 0.0	0.0 ± 0.0	0.0 ± 0.0	110.0 ± 9.0	3.0 – 2.6 ± 0.07
L1448IRS3A Outer Disk	3:25:36.50	30:45:21.84	84.4 ± 3.4	29.6 ± 0.5	0.75 × 0.39	136.5	0.73 ± 0.03	0.30 ± 0.02	134.1 ± 2.0	−99.0 ± −99.0	⋯

Note.

^aThe spectral index is calculated with respect to the 9.1 mm flux density given in Table 7 of Tychoniec et al. (2018b) under the assumption that F_ν ∝ ν^α. The range of spectral indices for most sources reflects the 9.1 mm flux density that is corrected and not corrected for free–free emission, respectively, in that study. The sources with only one spectral index listed could not be corrected for free–free emission. The formal uncertainly of the spectral index is quite low, ∼0.1, even if 10% systematic uncertainty in the flux density is assumed due to the large difference in wavelength from 1.3 to 9.1 mm. We calculated the uncertainty in the spectral index following Chiang et al. (2012, Equation A5).

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For the observed molecular lines, we made moment maps over the selected channel ranges where there is spectral line emission. These images are used to assess whether or not there is rotation around a particular multiple system and if there is an outflow.

3. Results

3.1. Dust Continuum Imaging

We detect a wide variety of structures toward the sample of 17 multiple protostar systems observed by ALMA. The dust continuum images of the Class 0 systems are shown in Figure 1 and the Class I systems in Figure 2. The ALMA observations had high enough angular resolution to resolve the companions in 15 out of the 17 systems. The systems with unresolved companions are Per-emb-2 and Per-emb-18, which have separations of 0 farcs 08 (∼24 au). For the 15 systems with resolvable companions, we detected all known components in 12. The companions were not clearly detected for Per-emb-40, Per-emb-55, and Per-emb-48. However, toward Per-emb-55, there is emission at the expected companion position at the 2σ level, and for Per-emb-40 there is extended emission toward the companion position that might reflect emission from the companion. Per-emb-48 was among the faintest sources detected by ALMA and self-calibration was not possible for this field. The increased noise from phase decorrelation may have prevented the companion from being detected, but we cannot rule out that its flux density is below our detection limit.

The Class 0 systems all have some extended emission surrounding the companions. This is demonstrated quantitatively in Table 1, where the total extended flux density (the sum of compact and extended emission) is greater than the flux density from the compact sources. Per-emb-2, Per-emb-18, Per-emb-17, L1448 IRS3B, SVS13A, L1448 IRS3C, and L1448 IRS2 all have clear emission surrounding and/or bridging the companion sources. The degree to which the total flux density is greater varies; sometimes the difference is as small as a few percent (e.g., Per-emb-17), but it can be greater than a factor of two (e.g., Per-emb-33). This emission could be described as a circumbinary disk and/or envelope.

The extended structure surrounding the companions is quite prominent in L1448 IRS3B, SVS13A, Per-emb-18, and Per-emb-2. The extended emission can be described as either spirals or streamers toward L1448 IRS3B, SVS13A, and Per-emb-2 (Figure 1). The spirals in L1448 IRS3B were previously reported in Tobin et al. (2016a). Finally, the extended structures surrounding Per-emb-2 and Per-emb-18 show evidence of being optically thick at 1.3 mm as evidenced by their emission peaks not being located toward what are assumed protostar positions, but off source, and having smooth surface brightness profiles across their central regions. In all these sources, the VLA imaging in Tobin et al. (2016b) detects the peak emission toward what are assumed to be the positions of the individual protostars, and the ALMA imaging with higher surface brightness sensitivity is able to detect much more extended dust emission than detected by the VLA. Furthermore, the brightest feature in Per-emb-18 is offset to the east; this was a weakly detected feature at 9 mm with the VLA (Tobin et al. 2016b).

Finally, there is less prominent extended emission toward NGC1333 IRAS2A and L1448 IRS2. NGC1333 IRAS2A has extended features at low surface brightness, and asymmetric extended emission to the east. Then the extended emission for L1448 IRS2 is also more diffuse and does not have as well-defined circumbinary structure like some of the others. However, the larger structure surrounding the protostars may be impacted by spatial filtering given that Tobin et al. (2015b) detected a surrounding structure on larger scales in lower-resolution data.

The Class I sources on the whole have little or no extended emission surrounding the components of these multiple systems. The images show this visually in Figure 2, and Table 1 shows that the flux densities for the Gaussian fits and the extended emission from a larger area encompassing the two protostars are comparable. However, some show resolved structure toward one component. Both NGC 1333 IRAS2B and L1448 IRS1 have at least one component that is dominant, with resolved structure in their dust emission. The other Class I binaries appear consistent with point sources, and the components have similar flux densities.

While the close companions were the primary targets, we detected several wide companions in the observed fields. In the field of SVS13A, we detected several additional sources. One source is RAC1999 VLA20 (Rodríguez et al. 1999), located northeast of SVS13A. This source was previously detected by the VLA, but has no counterpart at shorter wavelengths (i.e., mid-infrared and far-infrared) and has been hypothesized to be extragalactic (Tobin et al. 2016b), see Figure 3. The wider companion to SVS13A, often called VLA3 or SVS13A2, was detected (Figure 3), and appears marginally resolved. SVS13B is also detected and resolved as shown in Figure 3). VLA 8 mm imaging of SVS13B indicated that it has a small embedded disk (Segura-Cox et al. 2016), but a larger, resolved structure is detected in the ALMA 1.3 mm data. We also detected L1448 IRS3A in the field of L1448 IRS3B (Figure 4). The small-scale structure of L1448 IRS3A appears to be a resolved disk (Figure 4), consistent with the resolved emission detected in Tobin et al. (2015b). In the field of Per-emb-55 (Figure 5), we also detected Per-emb-8 which appears to have a large extended disk surrounding it (Figure 5). Per-emb-21 was detected in the field of Per-emb-18, but appears unresolved (Figure 6).

**Figure 3.** ALMA 1.3 mm images of the SVS13 (Per-emb-44) region. A wide view of the region encompassing all the sources within the primary beam is shown in the top left and dashed boxes mark the zoomed-in regions in the subsequent panels. The northeastern-most source (RAC1999 VLA20) is most likely a background extragalactic source from the VLA imaging results. The two sources south of SVS13A are SVS13A2 (VLA3) and SVS13B. We show the zoom-in on SVS13A2 in the top right panel, the zoom-in on SVS13B in the bottom left panel, and RAC1999 VLA20 in the bottom right panel. SVS13A2 and RAC1999 VLA20 do not appear highly resolved, while SVS13B is clearly resolved with respect to the beam.
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**Figure 4.** ALMA 1.3 mm images of the L1448 IRS3 (Per-emb-33) region encompassing all the sources within the primary beam; the dashed box marks the zoomed-in regions in the right panel. The observation was centered on L1448 IRS3B largest and brightest source near the bottom of the image. Nonetheless, L1448 IRS3A is detected (source in the middle left), and L1448 IRS3C (L1448 NW) is marginally detected in this image toward the upper right. L1448 IRS3C is also observed with a separate pointing given its large angular distance from field center. Right: ALMA 1.3 mm zoom-in on L1448 IRS3A showing the resolved disk structure toward this protostar.
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**Figure 5.** ALMA 1.3 mm image of the Per-emb-8/Per-emb-55 region the two sources within the primary beam; the dashed box marks the zoomed-in regions in the right panel. Per-emb-8 is located in the eastern portion of the image and Per-emb-55 is located in the western portion. Per-emb-8 is clearly resolved and Per-emb-55 is a binary, but only one component is strongly detected. Right: ALMA 1.3 mm image of Per-emb-8, an apparently single source with a surrounding disk. The central region of Per-emb-8 is very prominent with respect to the extended disk.
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**Figure 6.** ALMA 1.3 mm image of the Per-emb-18/Per-emb-21 region encompassing the two protostars within the primary beam; the dashed box marks the zoomed-in regions in the right panel. Per-emb-18 is located in the northern portion of the image and Per-emb-21 is located in the southern portion. Right: ALMA 1.3 mm image zooming-in on Per-emb-21, which appears unresolved.
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The resolved structures around Per-emb-8 and L1448 IRS3A are not ideal for fitting a single Gaussian, so we fit two Gaussians to these sources and list the inner and outer Gaussian fits separately, referring to them as the inner disk and outer disk, in addition to the fit of a single Gaussian in Tables 1 and 2.

3.2. Mass Estimates from the Dust Continuum

The mass of material found toward individual components of the multiple systems, as well as surrounding all components, can be estimated using the flux density of the dust continuum emission. Under the assumption that the dust emission is optically thin and isothermal, the dust mass can be calculated with the equation

$\begin{eqnarray}&&{M}_{\mathrm{dust}}=\displaystyle \frac{{D}^{2}{F}_{\nu }}{{\kappa }_{\nu }{B}_{\nu }({T}_{\mathrm{dust}})},\end{eqnarray} \tag{ 1 }$

where D is the distance (∼300 pc), F_ν is the observed flux density, B_ν is the Planck function, T_dust is the dust temperature, and κ_ν is the dust opacity at the observed wavelength. T_dust is assumed to be 30 K, consistent with temperature estimates on ∼100 au scales (Whitney et al. 2003) and κ_ν is 0.899, taken from Ossenkopf & Henning (1994, Table 1 column 5). We then multiply the resulting value of M_dust by 100, assuming the canonical dust-to-gas mass ratio of 1:100 (Bohlin et al. 1978). We caution that the masses calculated may have systematic uncertainty because opacity will hide some mass, the dust emission is not isothermal, and the dust opacity along with the dust-to-gas mass ratio are uncertain and may vary between components within the same system.

The results from the mass calculations are given in Table 2. A number of sources have total masses greater than 0.1 M_⊙. The mass calculations presented here tend to be about 20% higher than those reported for some of the same sources in Tobin et al. (2015b) from observations with CARMA (taking into account the different distances adopted), despite observing at the same wavelength, slightly lower resolution (for CARMA), and making the same assumptions for the mass calculation. The minor discrepancy likely results from a combination of greater absolute flux uncertainty in the CARMA data, better sensitivity to low-surface brightness emission from ALMA, better uv-coverage with ALMA, and less decoherence in the self-calibrated ALMA data.

Table 2. Compact Masses vs. Extended Masses

Source	Gaussian Mass	Extended Mass
	(M_⊙)	(M_⊙)
Per-emb-2	0.69 ± 0.01	0.68 ± 0.04
Per-emb-12	2.3 ± 0.04	3.9 ± 0.2
Per-emb-12-A	1.8 ± 0.04	⋯
Per-emb-12-B	0.53 ± 0.02	⋯
Per-emb-17	0.054 ± 0.001	0.066 ± 0.01
Per-emb-17-A	0.031 ± 0.0005	⋯
Per-emb-17-B	0.024 ± 0.0007	⋯
Per-emb-18	0.17 ± 0.002	0.18 ± 0.01
Per-emb-21	0.06 ± 0.0005	0.063 ± 0.005
Per-emb-22	0.053 ± 0.001	0.088 ± 0.014
Per-emb-22-A	0.041 ± 0.0009	⋯
Per-emb-22-B	0.012 ± 0.0002	⋯
Per-emb-27	0.24 ± 0.003	0.35 ± 0.03
Per-emb-27-A	0.204 ± 0.0024	⋯
Per-emb-27-B	0.032 ± 0.003	⋯
Per-emb-33	0.33 ± 0.003	0.89 ± 0.04
Per-emb-33-A	0.009 ± 0.001	⋯
Per-emb-33-B	0.04 ± 0.002	⋯
Per-emb-33-C	0.2 ± 0.002	⋯
L1448NW	0.08 ± 0.0009	0.1 ± 0.009
L1448NW-A	0.060 ± 0.0007	⋯
L1448NW-B	0.018 ± 0.0003	⋯
Per-emb-44	0.31 ± 0.002	0.48 ± 0.02
Per-emb-44-A	0.12 ± 0.001	⋯
Per-emb-44-B	0.19 ± 0.002	⋯
SVS13B	0.27 ± 0.003	0.33 ± 0.03
SVS13A2	0.017 ± 0.002	0.002 ± 0.005
RAC1999 VLA20	0.005 ± 0.001	0.003 ± 0.002
Per-emb-8	0.15 ± 0.003	0.15 ± 0.01
Per-emb-8 inner disk	0.11 ± 0.0009	⋯
Per-emb-8 outer disk	0.041 ± 0.003	⋯
Per-emb-35-A	0.027 ± 0.0005	⋯
Per-emb-35-B	0.020 ± 0.0005	⋯
Per-emb-36	0.17 ± 0.001	0.17 ± 0.005
Per-emb-36-A	0.15 ± 0.001	⋯
Per-emb-36-B	0.015 ± 0.0009	⋯
Per-emb-40-A	0.019 ± 0.0003	⋯
Per-emb-40-B	0.0017 ± 0.0002	⋯
Per-emb-48-A	0.005 ± 0.0005	0.005 ± 0.002
Per-emb-48-B	<0.0005	⋯
Per-emb-49-A	0.022 ± 0.0005	⋯
Per-emb-49-B	0.005 ± 0.0003	⋯
Per-emb-55-A	0.003 ± 0.0007	⋯
Per-emb-55-B	0.0007 ± 0.0003	⋯
L1448IRS1-A	0.088 ± 0.001	⋯
L1448IRS1-B	0.007 ± 0.0009	⋯
L1448IRS3A	0.12 ± 0.005	0.13 ± 0.01
L1448IRS3A inner disk	0.02 ± 0.001	⋯
L1448IRS3A outer disk	0.1 ± 0.004	⋯
EDJ2009-269-A	0.014 ± 0.0003	⋯
EDJ2009-269-B	0.01 ± 0.0003	⋯

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Most of the close companions have large flux density ratios between them, which could be interpreted as a large circumstellar mass ratio; this ratio offers no information as to the mass ratio between the stellar components and only reflects the compact circumstellar mass. Per-emb-17, EDJ2009-269, and Per-emb-35 have ratios closest to unity, and L1448 IRS1, L1448 NW, Per-emb-12, Per-emb-22, Per-emb-27, Per-emb-33, Per-emb-36, Per-emb-40, Per-emb-44, Per-emb-49, and Per-emb-55 all have circumstellar mass ratios that are less than 0.63. Both Class 0 and Class I systems span a range of circumstellar mass ratios. We note that several cases with large circumstellar mass ratios are in hierarchical and/or higher-order systems.

3.3. Spectral Indices

Because most of the companions are sufficiently resolved to enable their 1.3 mm flux densities to be measured individually, We were able to calculate the spectral index of the emission from 1.3 to 9.1 mm. We calculated the spectral index (α) using the equation

$\begin{eqnarray}&&\alpha =\displaystyle \frac{\mathrm{ln}({F}_{\nu ,1.3\mathrm{mm}})-\mathrm{ln}({F}_{\nu ,9.1\mathrm{mm}})}{\mathrm{ln}({\nu }_{1.3\mathrm{mm}})-\mathrm{ln}({\nu }_{9.1\mathrm{mm}})}\end{eqnarray} \tag{ 2 }$

and the associated uncertainty using the equation

$\begin{eqnarray}&&{\sigma }_{\alpha }^{2}={\left(\displaystyle \frac{1}{\mathrm{ln}({\nu }_{1.3\mathrm{mm}})-\mathrm{ln}({\nu }_{9.1\mathrm{mm}})}\right)}^{2}\left(\displaystyle \frac{{\sigma }_{\nu ,1.3\,\mathrm{mm}}^{2}}{{F}_{\nu ,1.3\,\mathrm{mm}}^{2}}+\displaystyle \frac{{\sigma }_{\nu ,9.1\,\mathrm{mm}}^{2}}{{F}_{\nu ,9.1\,\mathrm{mm}}^{2}}\right)\end{eqnarray} \tag{ 3 }$

following Chiang et al. (2012). We used the 9.1 mm flux densities from Tychoniec et al. (2018b) to calculate α with respect to the flux densities at 1.3 mm. Because Tychoniec et al. (2018b) attempted to remove the free–free emission that can contribute to the 9.1 mm flux densities, we calculated α with respect to both the corrected and uncorrected 9.1 mm flux densities. The spectral indices for the sources are given in Table 1. The statistical uncertainty on α is generally quite small, ∼0.1, due to the large difference in wavelength between the two bands. We assume a 10% uncertainty in the flux calibration for each flux density and add this in quadrature to the statistical uncertainty on the flux density. The uncertainty on α is generally smaller than the range of α from the corrected and uncorrected flux densities.

Most of the values for α are between 2 and 3; optically thin thermal dust emission, where the dust opacity spectral index κ_ν ∝ ν^β, is expected to have α = 2+β. Thus, for thermal dust emission alone (optically thin or partially optically thin), the spectral indices should be ≥2; the spectral index would be 2 for optically thick emission. Most of the values of α are >2 even when the 9.1 mm emission is not corrected for free–free emission, thus it is likely that the 9.1 mm mostly reflects dust emission as also found by Tychoniec et al. (2018b). There are some sources for which α ∼ 2; this could indicate that either the properties of the dust grains are such that β is quite small, the dust emission is becoming optically thick, or there is some amount of free–free emission to the 9.1 mm flux density that could not be removed. These three possibilities are difficult to disentangle with the data in hand, limiting the possible interpretations of these spectral index data. Therefore, the 1.3 mm emission presented here is most likely to reflect only dust emission while the 9.1 mm emission might reflect dust and free–free emission.

3.4. Molecular Line Kinematics

The dust continuum emission provides crucial evidence of the presence or absence of material surrounding the close binary/multiple stars, which will help with the interpretation of their formation mechanism. However, the continuum data alone are not enough to fully understand how these multiple systems formed. Therefore, we also observed five molecular tracers toward each system along with the dust continuum. We expect that C¹⁸O will generally trace the kinematics of the dense circum-multiple material because this line has tended to be the most reliable tracer of protostellar disks (Ohashi et al. 2014; Yen et al. 2014; Aso et al. 2015). Also, ¹³CO should generally trace the circum-multiple material (Takakuwa et al. 2012; Tobin et al. 2012), but we use ¹³CO with caution given that it can also trace outflowing gas. H₂CO is also expected to trace the circum-multiple environment because this is a typical high-density tracer (e.g., Mangum & Wootten 1993), and it has been found to trace the Class 0 disk and inner envelope of L1527 (Sakai et al. 2014). Similarly, SO has also been found to trace the disk toward some protostars (Sakai et al. 2014; Yen et al. 2014). However, there are instances where both H₂CO and SO have been shown to trace outflowing material, in addition to circumstellar/multiple material (e.g., Wakelam et al. 2005). Finally, ¹²CO is expected to trace the outflow toward most protostars in the sample given their youth, but ¹²CO could trace the disk kinematics of some of the more evolved multiple systems in our sample (e.g., Simon et al. 2000). The figures showing the molecular line detections toward circum-multiple and circumstellar structures are shown in Figures 7–22. The intervals over which the red and blueshifted emission are integrated are given in Table 3.

Table 3. Line Integration Intervals and Resultant Noise

Source	¹³CO Intervals	¹³CO rms	C¹⁸O Intervals	C¹⁸O rms	SO Intervals	SO rms	H₂CO Intervals	H₂CO rms	¹²CO Intervals	¹²CO rms
	(blue, red)	(blue, red)	(blue, red)	(blue, red)	(blue, red)	(blue, red)	(blue, red)	(blue, red)	(blue, red)	(blue, red)
	(km s⁻¹)	(K km s⁻¹)	(km s⁻¹)	(K km s⁻¹)	(km s⁻¹)	(K km s⁻¹)	(km s⁻¹)	(K km s⁻¹)	(km s⁻¹)	(K km s⁻¹)
Per-emb-2	5.0–6.0, 8.0–9.5	3.33, 3.95	7.5–9.0, 5.0–6.25	2.16, 1.65	⋯	⋯	⋯	⋯	1.5–6.5, 9.0–15.0	5.8, 6.15
Per-emb-12	1.25–6.5, 6.5–9.75	5.59, 6.65	4.25–6.75, 7.75–9.0	3.34, 2.46	1.0–7.0, 8.0–11.75	6.24, 4.99	3.75–7.0, 8.0–11.0	3.64, 3.51	−1.5–6.5, 10.0–15.5	7.21, 6.05
Per-emb-17	0.5–4.5, 6.5–9.75	6.08, 5.52	3.0–4.75, 5.5–7.75	2.19, 2.19	−1.5–5.0, 5.25–10.25	6.6, 5.8	0.0–5.0, 5.25–9.75	4.54, 4.32	0.0–4.0, 7.5–14.0	6.56, 5.26
Per-emb-18	4.75–7.5, 9.25–11.5	4.79, 4.54	4.75–7.25, 8.75–11.5	3.22, 3.37	4.5–8.0, 8.25–11.5	4.83, 4.67	4.5–7.5, 9.25,11.25	3.51, 3.08	3.0–5.0, 10.0 13.0	5.26, 4.41
Per-emb-21	⋯	⋯	⋯	⋯	⋯	⋯	⋯	⋯	⋯	⋯
Per-emb-22	2.25–3.75, 4.75–5.75	3.94, 3.33	2.75–4.0, 4.5 5.75	2.58, 2.58	1.75–5.75	5.21	3.0–4.0, 4.25–6.0	2.2, 2.79	−8.0–2.0, 7.5–16.0	8.1, 7.5
Per-emb-27	0.25–6.25, 9.5–10.5	7.03, 3.14	5.75–7.25, 7.5–9.0	2.58, 2.58	0.25–7.0, 7.25–11.0	7.0, 5.29	3.0–6.75, 7.75–10.0	3.91, 3.09	−28.5–6.0, 9.5–20.0	13.56, 13.58
Per-emb-33	1.25-4, 5.5–7.0	4.99, 3.20	1.25–4.0, 5.5–7.0	2.25, 1.65	1.25–4.5, 4.5–8.0	4.74, 4.91	2.75–4.0, 5.25–6.25	2.25, 2.05	−5.5–1.5, 4.5–8.0	6.88, 5.02
L1448NW	0.75–3.5, 5.25–7.75	4.95, 4.74	1.0–3.75, 4.0, 7.25	3.66, 3.39	2.25–5.0	4.41	1.5–6.25	4.12	−9.0–3.5, 6.0–16.0	8.4, 7.54
Per-emb-44	6.75–7.5, 9.0, 11.0	2.78, 4.17	6.0–8.5, 8.75, 10.0	3.63, 2.34	4.75–8.5, 8.75–11.75	6.36, 5.73	4.25–8.5, 8.75–11.75	3.63, 2.34	−25.0–0.0, 10.0–20.0	4.42, 4.22
Per-emb-8	5.75–8.75, 12.0–14.5	5.09, 4.68	6.0–9.25, 11.75–14.0	3.57, 3.02	7.5–10.0, 10.75–12.0	4.04, 2.99	6.75–10.25, 11.0–13.5	3.42, 2.93	4.5–8.5, 13.0–16.5	4.89, 4.61
Per-emb-55	⋯	⋯	⋯	⋯	⋯	⋯	⋯	⋯	4.5–8.5, 11.5–16.0	4.90, 5.16
Per-emb-35	5.75–6.25, 6.5–7.25	2.52, 2.91	5.25–8.0	4.31	5.5–7.25, 7.5–9.5	3.35, 3.55	5.5–7.5, 7.75–10.5	2.87, 3.63	2.0–7.5, 10.0–14.0	6.0, 5.2
Per-emb-36	3.5–5.75, 9.0–11.0	4.68, 4.17	4.5–6.5	3.14	4.5–7.0, 7.25–9.0	4.29, 3.65	4.0–6.25	2.87	−3.0–5.0, 9.5–16.0	7.81, 7.1
Per-emb-40	5.0–6.5	3.91, 4.18	5.0–7.75	3.38	⋯	⋯	⋯	⋯	−3.5–5.0, 8.5–13.5	8.03, 6.28
Per-emb-48	⋯	⋯	⋯	⋯	⋯	⋯	⋯	⋯	⋯	⋯
Per-emb-49	⋯	⋯	⋯	⋯	⋯	⋯	⋯	⋯	⋯	⋯
L1448IRS1	⋯	⋯	⋯	⋯	⋯	⋯	⋯	⋯	0.0–2.0, 5.0–7.0	3.65, 3.65
L1448IRS3A	2.5–3.75, 6.25–9.0	2.96, 4.19	2.25–3.75, 6.25–9.0	2.42, 3.18	3.0–4.25, 4.25–8.0	3.1, 5.07	3.0–3.75, 5.25–7.0	1.83, 2.6	−1.0–1.5, 6.5–11.0	4.35, 5.61
EDJ2009-269	⋯	⋯	⋯	⋯	⋯	⋯	⋯	⋯	5.5–8.0	5.04

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Due to the short observations toward each source, the S/N is not very high in each channel for these high angular resolution spectral line maps. Thus, we examine the integrated intensity maps constructed for emission red- and blueshifted with respect to the line center. The detection rates for the lines vary from source to source; some only have one line detected, whereas others have detections in all five lines. We examine the integrated intensity maps of each molecule, relative to the outflow directions and the extended continuum structure. The larger-scale outflows have previous characterization of their orientations, when possible, from Stephens et al. (2017); therefore we only show the outflow maps in Appendix A.

3.4.1. Characterization of Kinematics in Multiples

We use the red- and blueshifted integrated intensity maps of each molecule to determine whether or not the velocity patterns observed toward the multiple systems reflect rotation, outflow, or an indistinct process. We do this by comparing the orientation of the velocity gradient relative to the outflow direction (and extended continuum if possible). If there is a velocity gradient in the circum-multiple structure that is within 30° of orthogonal to the outflow direction, we classify it as being consistent with disk rotation. To quantify the relative orientations of the outflow and velocity gradient directions, we measure the position angle of both and then take the difference. These values are given in Table 4. To measure the outflow position angle, we draw a line bisecting the outflow and measure its position angle. We follow the same procedure for the velocity gradients, but measure this from the emission peaks in the blue- and redshifted integrated intensity maps tracing rotation, generally C¹⁸O. The position angles have uncertainty and we estimate this to be ±10°.

Table 4. Molecules and Rotation Characteristics

Source	Class	Separation	Outflow PA^a	Gradient PA^b	¹²CO	¹³CO	C¹⁸O	H₂CO	SO	Rotation versus Outflow	Disk Frag. Possible?
		('', au)	(°)	(°)
Per-emb-2	0	0.080, 24.0	128	232 (104)	Outflow	Rotation	Rotation	⋯	⋯	Perpendicular	Yes
Per-emb-12	0	1.830, 548.9	175, 205	⋯	Outflow	Indistinct	Indistinct	Indistinct	Indistinct	Indistinct	No
Per-emb-17	0	0.278, 83.3	240	180/133 (60/107)	Outflow	Rotation	Rotation	Outflow+Rotation	Outflow+Rotation	Perpendicular	Yes
Per-emb-18	0	0.085, 25.6	169	78 (91)	Outflow	Rotation	Rotation	Rotation	Rotation	Perpendicular	Yes
Per-emb-22	0	0.751, 225.4	306	246 (60)	Outflow	Rotation	Rotation	Indistinct	Indistinct	Perpendicular	Yes
Per-emb-27	0	0.620, 186.0	204, 285	⋯	Outflow	Indistinct	Indistinct	Outflow/Indistinct	Outflow	Indistinct	No
Per-emb-33	0	0.264, 79.2	301, 284	212 (89)	Outflow	Rotation	Rotation	Rotation	Outflow?	Perpendicular	Yes
	0	0.795, 238.4
L1448NW	0	0.251, 75.3	305	224 (81)	Outflow	Rotation	Rotation	Indistinct	Indistinct	Perpendicular	Yes
Per-emb-44	0/I	0.300, 90.0	147	220 (73)	Outflow	Rotation	Rotation	Rotation?	Rotation?	Perpendicular	Yes
Per-emb-35	I	1.908, 572.3 292, 279	⋯	Outflow	Outflow?	Indistinct	Indistinct	Indistinct	Indistinct	No?
Per-emb-36	I	0.311, 93.4	24	⋯	Outflow	Indistinct	Indistinct	Rotation?	Outflow?	Indistinct	No?
Per-emb-40	I	0.391, 117.4	123?	⋯	Outflow/Rotation	Indistinct	Indistinct	⋯	⋯	⋯	⋯
Per-emb-48	I	0.346, 103.7	⋯	⋯	Indistinct	⋯	⋯	⋯	⋯	⋯	⋯
Per-emb-49	I	0.313, 93.8	⋯	⋯	⋯	⋯	⋯	⋯	⋯	⋯	⋯
Per-emb-55	I	0.618, 185.3	121?	⋯	Outflow	⋯	⋯	⋯	⋯	⋯	⋯
L1448IRS1	I	1.424, 427.0	113	212 (99)	Rotation	⋯	⋯	⋯	⋯	Perpendicular	Yes?
EDJ2009-269	II	0.524, 157.3	⋯	⋯	Indistinct	⋯	⋯	⋯	⋯	⋯	⋯
Per-emb-35-A	I		292	201 (91)	Outflow	Outflow?	Indistinct	Rotation?	Indistinct	Perpendicular?
Per-emb-35-B	I		279	233 (46)	Outflow	Outflow?	Indistinct	Rotation?	Indistinct	Perpendicular?
L1448IRS3A	I		38	313 (85)	Outflow?	Rotation	Rotation	Rotation	Rotation	Perpendicular
Per-emb-8	0		314	44 (90)	Outflow/Rotation	Rotation	Rotation	Rotation	Rotation	Perpendicular
Per-emb-21	0		80?	⋯	Outflow	⋯	⋯	⋯	⋯	⋯
SVS13B	0		160	⋯	Outflow	⋯	⋯	⋯	⋯	⋯

Notes. This table examines what kinematics are being displayed by each molecular line. The ellipsis symbols (...) correspond to non-detections. The question marks following a descriptor indicate that the line might display a certain kinematic signature, but that ambiguity remains. We expect the outflow position angle (PA) and gradient PA to have a typical uncertainty of 10°.

^aThis is measured from the ¹²CO data presented in this paper except for Per-emb-36 and L1448 IRS1. To measure the outflow PA, we draw a line that best bisects the outflow cavity. We adopt the outflow PA value from Plunkett et al. (2013) for Per-emb-36 because of how wide the outflow is. For L1448 IRS1, ¹²CO is tracing rotation, thus we use the scattered light nebula in the image from Foster & Goodman (2006). ^bThe PA of the velocity gradient is measured by drawing a line from the peaks of the blue- and redshifted integrated intensity maps, generally C¹⁸O.

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Molecular line emission is detected with sufficiently high S/N in 12 systems to characterize the kinematics as rotation or not. This classification is most robust for the nine Class 0 systems with the strongest detections and seven of these Class 0 systems have velocity gradients that indicate rotation in the circum-multiple gas. If we include those with tentative classifications, we find that eight out of 12 systems are consistent with the circum-multiple emission tracing rotation orthogonal to the outflow direction. The systems classified as indistinct generally have disorganized kinematic structure on the scales and at the sensitivity we observed. Also, certain molecules may better trace different kinematic components of protostellar systems (e.g., Sakai et al. 2014). The detection and characterization of the molecular line data are summarized in Table 4. Each source is discussed in more detail in Section 4.1 with respect to the data presented in this paper and to interpret previous results in the context of the new ALMA data.

3.4.2. Detections of Individual Rotating Disks

In addition to the kinematic structure of circum-multiple material, we were also able to examine the kinematic structure of both some individual circumstellar disks either as part of a close multiple system or as a wide companion. Per-emb-8 shows the clearest example of rotation in all observed molecular lines; ¹²CO interestingly shows a combination of rotation and outflow. L1448 IRS3A also shows rotation in all molecular lines and is oriented orthogonal to the position angle of L1448 IRS3B. The individual disks toward Per-emb-35-A and -B appear to have molecular line emission as well, but this is not well-resolved spatially. Only H₂CO appears to have a clear indication of rotation. The characterization of the observed kinematics from all lines is given in Table 4.

4. Discussion

The ALMA 1.3 mm continuum and molecular line emission observed toward the close multiple systems in Perseus dramatically increase our knowledge of the spatial and kinematic structure of dense gas on a few hundred au scales around these multiple protostar systems. The ALMA dust continuum traces the more tenuous emission surrounding the companions as compared to previous VLA observations. While the extended emission is more diffuse than the compact sources detected by both the VLA and ALMA, this emission can account for a significant fraction of the mass in the systems. The molecular maps show the gas motion, revealing how well (or not) the motion of the extended dense gas surrounding the protostar(s) is organized. The molecular line maps provide evidence for rotation in several systems, but not universally, as some systems have unclear kinematic structure with the current S/N of the ALMA data. Nevertheless, these ALMA data now enable us to better interpret how multiple star systems form on a case-by-case basis and as an ensemble. We consider each system individually and summarize our findings for the sample as a whole.

4.1. Individual Systems

4.1.1. Class 0 and 0/I Systems

Per-emb-2. The continuum structure is extended with a very smooth surface brightness distribution, in contrast to the VLA data with a central intensity peak toward the protostars, forming a binary system with a separation of ∼24 au. Only C¹⁸O and ¹³CO are detected toward this system, and we detect a clear velocity gradient orthogonal to the outflow direction, as shown in Figure 7. The C¹⁸O emission is not fully coincident with the continuum image, possibly due to dust opacity and/or spatial filtering of low-velocity emission near line-center. Nevertheless, the C¹⁸O and ¹³CO kinematics still exhibit evidence for rotation on the scale of the circumbinary continuum.

**Figure 7.** Integrated intensity maps of ¹³CO (left panel) and C¹⁸O (right panel) toward Per-emb-2. The integrated intensity maps are displayed as red and blue contours corresponding to the integrated intensity of line emission red- and blueshifted with respect to the system velocity. The contours are overlaid on the 1.3 mm continuum image. The line emission shows evidence for a velocity gradient consistent with rotation. The redshifted contours start at (4, 3)σ and increase in (1, 1)σ increments, and the blueshifted contours start at (3, 3)σ and increase in (1, 1)σ increments. The values inside the parentheses in the previous sentence correspond to the ¹³CO, C¹⁸O, SO, and H₂CO integrated intensity maps, respectively. The values for σ_red and σ_blue and velocity ranges over which the line emission was summed can be found in Table 3. The beam in the images is approximately 036 × 026.
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Per-emb-12. Also known as NGC 1333 IRAS4A, this system is the most massive overall in terms of dust emission out to very large scales. It was also known to be a binary prior to the VANDAM survey with a separation of ∼549 au (Looney et al. 2000). The dust emission toward IRAS4A has significant opacity, which prevents tracing kinematics toward the protostars; only some emission around the protostars in C¹⁸O and ¹³CO is revealed in Figure 8.

**Figure 8.** Integrated intensity maps of ¹³CO (top left panel), C¹⁸O (top right panel), SO (bottom left panel), and H₂CO (bottom right panel) toward Per-emb-12 (NGC 1333 IRAS4A). The integrated intensity maps are displayed as red and blue contours corresponding to the integrated intensity of line emission red- and blueshifted with respect to the system velocity. The contours are overlaid on the 1.3 mm continuum image. The line emission shows evidence for a velocity gradient consistent with rotation. The redshifted contours start at (3, 3, 4, 3)σ and increase in (3, 1, 2, 3)σ increments, and the blueshifted contours start at (3, 3, 4, 3)σ and increase in (3, 1, 2, 3)σ increments. The values inside the parentheses in the previous sentence correspond to the ¹³CO, C¹⁸O, SO, and H₂CO integrated intensity maps, respectively. The values for σ_red and σ_blue and velocity ranges over which the line emission was summed can be found in Table 3. The beam in the images is approximately 036 × 026.
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Per-emb-17. There were detections in multiple lines toward Per-emb-17, and the material surrounding the ∼83 au binary system forms a flattened structure in the dust continuum that is similarly flat in ¹³CO and C¹⁸O. The molecular line emission exhibits velocity gradients orthogonal to the outflow direction; see Figure 9. The rotation center of the system is not clearly centered on a particular component, so it is unclear which protostar is more massive. H₂CO and SO, however, seem to be more associated with the outflow in this source. SO in particular is strong toward the B component and extended along the outflow.

**Figure 9.** Integrated intensity maps of ¹³CO (top left panel), C¹⁸O (top right panel), SO (bottom left panel), and H₂CO (bottom right panel) toward Per-emb-17. The integrated intensity maps are displayed as red and blue contours corresponding to the integrated intensity of line emission red- and blueshifted with respect to the system velocity. The contours are overlaid on the 1.3 mm continuum image. The line emission shows evidence for a velocity gradient consistent with rotation. The redshifted contours start at (3, 4, 3, 3)σ and increase in (2, 1, 2, 3)σ increments, and the blueshifted contours start at (3, 3, 3, 3)σ and increase in (2, 1, 2, 3)σ increments. The values inside the parentheses in the previous sentence correspond to the ¹³CO, C¹⁸O, SO, and H₂CO integrated intensity maps, respectively. The values for σ_red and σ_blue and velocity ranges over which the line emission was summed can be found in Table 3. The beam in the images is approximately 036 × 026.
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Per-emb-18. This source has a flattened structure in the dust continuum surrounding the companions separated by ∼26 au and detections in multiple lines. The continuum might be tracing dust emission with high optical depth, given that the emission is not very peaked toward the main pair of protostars. ¹³CO, C¹⁸O, SO, and H₂CO trace velocity gradients orthogonal to the outflow direction; see Figure 10. The fact that all molecules have gradients in the same direction lends confidence to interpreting the kinematics as rotation.

**Figure 10.** Integrated intensity maps of ¹³CO (top left panel), C¹⁸O (top right panel), SO (bottom left panel), and H₂CO (bottom right panel) toward Per-emb-18. The integrated intensity maps are displayed as red and blue contours corresponding to the integrated intensity of line emission red- and blueshifted with respect to the system velocity. The contours are overlaid on the 1.3 mm continuum image. The line emission shows evidence for a velocity gradient consistent with rotation. The redshifted contours start at (3, 3, 3, 3)σ and increase in (2, 1, 2, 2)σ increments, and the blueshifted contours start at (3, 4, 3, 3)σ and increase in (2, 1, 2, 2)σ increments. The values inside the parentheses in the previous sentence correspond to the ¹³CO, C¹⁸O, SO, and H₂CO integrated intensity maps, respectively. The values for σ_red and σ_blue and velocity ranges over which the line emission was summed can be found in Table 3. The beam in the images is approximately 036 × 026.
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Per-emb-22. Also known as L1448 IRS2, this binary system has a separation of ∼225 au. Continuum emission was mainly detected toward individual components, with some emission in between, exhibiting an apparent "bridge" between them. The molecular lines have evidence of velocity gradients orthogonal to the outflow (Figure 11), but much of the molecular line emission appears to be filtered out and/or have very low surface brightness. Tobin et al. (2015b) previously found evidence for velocity gradients/rotation using combined CARMA and SMA data that had lower resolution, in addition to a larger circumbinary structure.

**Figure 11.** Integrated intensity maps of ¹³CO (top left panel), C¹⁸O (top right panel), SO (bottom left panel), and H₂CO (bottom right panel) toward Per-emb-22 (L1448 IRS2). The integrated intensity maps are displayed as red and blue contours corresponding to the integrated intensity of line emission red- and blueshifted with respect to the system velocity. The green contours for the SO emission correspond to emission integrated across the entire line because there was not a clear velocity gradient in this molecule due to the low S/N. The contours are overlaid on the 1.3 mm continuum image. The line emission shows evidence for a velocity gradient consistent with rotation. The redshifted contours start at (3, 3, 3, 4)σ and increase in (3, 3, 3, 2)σ increments, and the blueshifted contours start at (3, 3, 3, 4)σ and increase in (3, 3, 3, 2)σ increments. The values inside the parentheses in the previous sentence correspond to the ¹³CO, C¹⁸O, SO, and H₂CO integrated intensity maps, respectively. The values for σ_red and σ_blue and velocity ranges over which the line emission was summed can be found in Table 3. The beam in the images is approximately 036 × 026.
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Per-emb-27. Also known as NGC 1333 IRAS2A, this system does not have much material surrounding the ∼186 au binary, but there is some extended material associated with the brighter source. There is weak evidence for a velocity gradient in C¹⁸O on scales larger than the binary system. SO is associated with the outflow, but H₂CO is compact toward the brighter component with a small extension toward the companion; see Figure 12.

**Figure 12.** Integrated intensity maps of ¹³CO (top left panel), C¹⁸O (top right panel), SO (bottom left panel), and H₂CO (bottom right panel) toward Per-emb-27 (NGC 1333 IRAS2A). The integrated intensity maps are displayed as red and blue contours corresponding to the integrated intensity of line emission red- and blueshifted with respect to the system velocity. The contours are overlaid on the 1.3 mm continuum image. The line emission shows evidence for a velocity gradient consistent with rotation. The redshifted contours start at (3, 3, 3, 3)σ and increase in (3, 1, 2, 3)σ increments, and the blueshifted contours start at (3, 3, 3, 3)σ and increase in (3, 1, 2, 3)σ increments. The values inside the parentheses in the previous sentence correspond to the ¹³CO, C¹⁸O, SO, and H₂CO integrated intensity maps, respectively. The values for σ_red and σ_blue and velocity ranges over which the line emission was summed can be found in Table 3. The beam in the images is approximately 036 × 026.
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Per-emb-33/L1448 IRS3B. This source was shown in Tobin et al. (2016a) to have clear velocity gradients associated with its extended circumbinary material in ¹³CO, C¹⁸O, and H₂CO. The separations of the companions from Per-emb-33-A are ∼79 and ∼238 au. Despite being published previously, we show this source for the sake of comparison to the full sample; see Figure 13.

**Figure 13.** Integrated intensity maps of ¹³CO (top left panel), C¹⁸O (top right panel), SO (bottom left panel), and H₂CO (bottom right panel) toward Per-emb-33 (L1448 IRS3B). The integrated intensity maps are displayed as red and blue contours corresponding to the integrated intensity of line emission red- and blueshifted with respect to the system velocity. The contours are overlaid on the 1.3 mm continuum image. The line emission shows evidence for a velocity gradient consistent with rotation. The redshifted contours start at (4, 4, 3, 3)σ and increase in (1, 1, 2, 2)σ increments, and the blueshifted contours start at (4, 4, 3, 3)σ and increase in (1, 1, 2, 2)σ increments. The values inside the parentheses in the previous sentence correspond to the ¹³CO, C¹⁸O, SO, and H₂CO integrated intensity maps, respectively. The values for σ_red and σ_blue and velocity ranges over which the line emission was summed can be found in Table 3. The beam in the images is approximately 036 × 026.
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L1448 IRS3C. The two companions in this system are separated by just ∼75 au, and they appear blended in the image. However, the emission toward each component is measured using a two-component Gaussian fit. There is low-level dust continuum emission surrounding the two blended protostars, and the circumbinary material also shows a velocity gradient in C¹⁸O orthogonal to the outflow; see Figure 14.

**Figure 14.** Integrated intensity maps of ¹³CO (top left panel), C¹⁸O (top right panel), SO (bottom left panel), and H₂CO (bottom right panel) toward L1448 IRS3C (L1448 NW). The integrated intensity maps are displayed as red and blue contours corresponding to the integrated intensity of line emission red- and blueshifted with respect to the system velocity. The contours are overlaid on the 1.3 mm continuum image. The line emission shows evidence for a velocity gradient consistent with rotation. The redshifted contours start at (3, 4)σ and increase in (2, 1)σ increments, and the blueshifted contours start at (3, 3)σ and increase in (2, 1)σ increments. The contours for SO and H₂CO start (3, 3)σ and increase in (2, 3)σ increments, respectively, for a single interval. The values for σ, σ_red and σ_blue and velocity ranges over which the line emission was summed can be found in Table 3. The beam in the images is approximately 036 × 026.
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Per-emb-44. Also known as SVS 13A, this binary system has a companion separation of ∼90 au and is one of the most prominent protostars within NGC1333. The 1.3 mm continuum also shows a one-armed spiral pattern, almost as prominent as that of L1448 IRS3B. Both ¹³CO and C¹⁸O are detected toward this protostar and exhibit velocity gradients that are at least partially orthogonal to the outflow; see Figure 15. This source might be viewed more face-on making rotation harder to distinguish in this complex environment. The C¹⁸O emission also follows the spiral through multiple velocity channels. Furthermore, both SO and H₂CO are detected at smaller radii near the close pair and may also show a velocity gradient, but this is uncertain due to blending. This protostar was previously known to have a close companion (Anglada et al. 2004), and the original study indicated that the disk around the western companion (Per-emb-44-B; VLA 4A) was not very prominent given its low flux density relative to the eastern component (Per-emb-44-A; VLA 4B) at a wavelength of 7 mm. However, their masses (flux densities) differ by only a factor of 1.6 at 1.3 mm, which may have more to do with either the amount of large grains and/or variable free–free emission because Anglada et al. (2004) used data at a wavelength of 7 mm from the VLA. However, Tychoniec et al. (2018b) also finds that the 9 mm flux densities (with estimated free–free contribution removed) have a flux ratio of 2.8, and the spectral indices are quite different, ∼2.3 for the eastern component and ∼3 for the western component. This could indicate more large grains in the eastern source to explain the increased flux density, or the 1.3 mm is optically thick and the 8 mm emission is less opaque, enabling more dust emission to be detected.

**Figure 15.** Integrated intensity maps of ¹³CO (top left panel), C¹⁸O (top right panel), SO (bottom left panel), and H₂CO (bottom right panel) toward Per-emb-44 (SVS13A). The integrated intensity maps are displayed as red and blue contours corresponding to the integrated intensity of line emission red- and blueshifted with respect to the system velocity. The contours are overlaid on the 1.3 mm continuum image. The line emission shows evidence for a velocity gradient consistent with rotation. The redshifted contours start at (4, 3, 5, 5)σ and increase in (2, 1, 5, 5)σ increments, and the blueshifted contours start at (4, 3, 5, 5)σ and increase in (2, 1, 5, 5)σ increments. The values inside the parentheses in the previous sentence correspond to the ¹³CO, C¹⁸O, SO, and H₂CO integrated intensity maps, respectively. The values for σ_red and σ_blue and velocity ranges over which the line emission was summed can be found in Table 3. The beam in the images is approximately 036 × 026.
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4.1.2. Class I and II Systems

Per-emb-36 (NGC 1333 IRAS2B). The two sources detected toward Per-emb-36 are marginally resolved in the ALMA 1.3 mm data with a separation of ∼93 au, and have multiple detected emission lines (Figure 16). However, the line emission does not obviously trace velocity gradients across the source, and the emission is mainly found west of the protostars rather than around them.

**Figure 16.** Integrated intensity maps of ¹³CO (top left panel), C¹⁸O (top right panel), SO (bottom left panel), and H₂CO (bottom right panel) toward Per-emb-36 (NGC 1333 IRAS2B). The integrated intensity maps are displayed as red and blue contours corresponding to the integrated intensity of line emission red- and blueshifted with respect to the system velocity. The contours are overlaid on the 1.3 mm continuum image. The line emission shows evidence for a velocity gradient consistent with rotation. The redshifted contours start at (4, 3)σ and increase in (2, 3)σ increments for ¹³CO and SO, respectively, and the blueshifted contours start at (3, 3, 3)σ and increase in (2, 3, 2)σ increments for ¹³CO, SO, and H₂CO, respectively. The line-center contours for C¹⁸O start at 3σ and increase by 1σ increments. The values for σ_red and σ_blue and velocity ranges over which the line emission was summed can be found in Table 3. The beam in the images is approximately 036 × 026.
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Per-emb-35. These protostars are separated by ∼572 au, and they are the source of an S-shaped outflow detected by Spitzer (Gutermuth et al. 2008). These sources do not have obvious dust continuum surrounding the compact continuum sources, but we cannot rule out diffuse circumbinary emission below our detection limits. H₂CO and SO are strongly detected toward both of the protostars and may show evidence of velocity gradients orthogonal to their outflow directions (Figure 17). Furthermore, there appears to be emission in these molecules between the two protostars. Both protostars are driving outflows, and they appear parallel as observed in the plane of the sky. ¹³CO is detected toward both sources, but it is not obviously tracing rotation, and C¹⁸O is only weakly detected; see Figure 17.

**Figure 17.** Integrated intensity maps of ¹³CO (top left panel), C¹⁸O (top right panel), SO (bottom left panel), and H₂CO (bottom right panel) toward Per-emb-35 (NGC 1333 IRAS1). The integrated intensity maps are displayed as red and blue contours corresponding to the integrated intensity of line emission red- and blueshifted with respect to the system velocity. The contours are overlaid on the 1.3 mm continuum image. The line emission shows evidence for a velocity gradient consistent with rotation. The redshifted contours start at (3, 3, 3, 3)σ and increase in (2, 1, 2, 2)σ increments and the blueshifted contours start at (3, 3, 3, 3)σ and increase in (2, 1, 2, 2)σ increments, respectively. The values inside the parentheses in the previous sentence correspond to the ¹³CO, C¹⁸O, SO, and H₂CO integrated intensity maps, respectively. The values for σ_red and σ_blue and velocity ranges over which the line emission was summed can be found in Table 3. The beam in the images is approximately 036 × 026.
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Per-emb-40. Only the primary component of this system is well detected, but the companion is tentatively detected. The molecules ¹³CO and C¹⁸O are only weakly detected (Figure 18). ¹³CO does have blue- and redshifted components, but not obviously tracing rotation.

**Figure 18.** Integrated intensity maps of ¹³CO (top left panel), C¹⁸O (top right panel), SO (bottom left panel), and H₂CO (bottom right panel) toward Per-emb-40 (B1-a). The integrated intensity maps are displayed as red and blue contours corresponding to the integrated intensity of line emission red- and blueshifted with respect to the system velocity. The contours are overlaid on the 1.3 mm continuum image. The line emission shows evidence for a velocity gradient consistent with rotation. The redshifted contours start at 3σ and increase in 2σ increments, and the blueshifted contours start at 3σ and increase in 2σ increments. For C¹⁸O only the line center velocities are plotted, starting at 4σ and increasing by 1σ intervals. The values for σ_red and σ_blue and velocity ranges over which the line emission was summed can be found in Table 3. The beam in the images is approximately 036 × 026.
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L1448 IRS1. This protostar appears to be the most evolved source in L1448, and it is detectable at visible wavelengths. The separation of the two protostars is ∼427 au. Unlike the other protostellar systems, the rotation in L1448 IRS1 is traced by ¹²CO that is extended toward the companion; see Figure 19. In this particular case, ¹²CO, is not tracing outflowing emission but rather the disk because of its advanced evolutionary state. We know this because the scattered light nebula from this source (Foster & Goodman 2006), which is an alternative indicator of the outflow direction, is oriented orthogonal to the velocity gradient in ¹²CO.

**Figure 19.** Integrated intensity maps of ¹²CO toward L1448 IRS1. The integrated intensity maps are displayed as red and blue contours corresponding to the integrated intensity of line emission red- and blueshifted with respect to the system velocity. The contours are overlaid on the 1.3 mm continuum image. The line emission shows evidence for a velocity gradient consistent with rotation. The redshifted contours start at 3σ and increase in 3σ increments, and the blueshifted contours start at 3σ and increase in 3σ increments. The values for σ_red and σ_blue and velocity ranges over which the line emission was summed can be found in Table 3. The beam in the images is approximately 036 × 026.
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EDJ2009-269. This system is a likely Class II binary in Perseus, separated by ∼157 au. Its only detected emission line is ¹²CO toward one of the protostars (Figure 20). In this more evolved system, the ¹²CO also appears to trace kinematics of gas toward the individual circumstellar disks rather than outflowing material.

**Figure 20.** Integrated intensity maps of ¹²CO toward EDJ2009-269. The integrated intensity maps are displayed as red and blue contours corresponding to the integrated intensity of line emission red- and blueshifted with respect to the system velocity. The contours are overlaid on the 1.3 mm continuum image. The line emission shows evidence for a velocity gradient consistent with rotation. The integrated intensity contours start at 3σ and increase in 3σ increments. The values for σ and the velocity range over which the line emission was summed can be found in Table 3. The beam in the images is approximately 036 × 026.
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Per-emb-48, Per-emb-49, and Per-emb-55. These protostars do not show strong evidence for molecular line emission surrounding the companion stars, thus we cannot examine whether or not there are velocity gradients in the circum-multiple emission. These are all Class I systems and may simply have a much smaller reservoir of circumbinary material, and/or the molecular line emission could not be detected with our uv-coverage and/or sensitivity. Per-emb-55, however, does appear to have an outflow associated with it in ¹²CO emission; see the Appendix, Figure 24.

4.1.3. Wide Companions with Disks

L1448 IRS3A. This is the wide companion to L1448 IRS3B, separated by 7 farcs 3 (∼2195 au), and has extended dust emission with disk-like morphology. The protostar exhibits a weak outflow in ¹²CO, but rotation in all the other observed lines; see Figure 21.

**Figure 21.** Integrated intensity maps of ¹³CO (top left panel), C¹⁸O (top right panel), SO (bottom left panel), and H₂CO (bottom right panel) toward L1448 IRS3A. The integrated intensity maps are displayed as red and blue contours corresponding to the integrated intensity of line emission red- and blueshifted with respect to the system velocity. The contours are overlaid on the 1.3 mm continuum image. The line emission shows evidence for a velocity gradient consistent with rotation. The redshifted contours start at (3, 4, 3, 3)σ and increase in (2, 1, 2, 3)σ increments, and the blueshifted contours start at (3, 3, 3, 3)σ and increase in (2, 1, 2, 3)σ increments. The values inside the parentheses in the previous sentence correspond to the ¹³CO, C¹⁸O, SO, and H₂CO integrated intensity maps, respectively. The values for σ_red and σ_blue and velocity ranges over which the line emission was summed can be found in Table 3. The beam in the images is approximately 036 × 026.
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Per-emb-8. This protostar is clearly detected with extended dust emission that shows signatures of rotation from the molecular line emission (Figure 22). It is separated from Per-emb-55 by ∼2867 au. Per-emb-8 was previously identified as a disk candidate by Segura-Cox et al. (2016), but with a very different position angle than is observed with ALMA. Since this source appears to have strong, extended free–free emission (Tychoniec et al. 2018a), the geometric parameters from the VLA data may not solely trace the dusty disk.

**Figure 22.** Integrated intensity maps of ¹³CO (top left panel), C¹⁸O (top right panel), SO (bottom left panel), and H₂CO (bottom right panel) toward Per-emb-8. The integrated intensity maps are displayed as red and blue contours corresponding to the integrated intensity of line emission red- and blueshifted with respect to the system velocity. The contours are overlaid on the 1.3 mm continuum image. The line emission shows evidence for a velocity gradient consistent with rotation. The redshifted contours start at (3, 4, 3, 3)σ and increase in (2, 1, 2, 2)σ increments, and the blueshifted contours start at (3, 3, 3, 3)σ and increase in (2, 1, 2, 2)σ increments. The values inside the parentheses in the previous sentence correspond to the ¹³CO, C¹⁸O, SO, and H₂CO integrated intensity maps, respectively. The values for σ_red and σ_blue and velocity ranges over which the line emission was summed can be found in Table 3. The beam in the images is approximately 036 × 026.
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4.2. The Origin of Close Protostellar Multiples

The ALMA continuum and molecular line maps are key to determining the most likely origin of multiple star systems. Systems that have formed via disk fragmentation are expected to have circum-multiple dust emission and this circum-multiple material should be rotating in an organized manner. Furthermore, if fragmentation occurred in a disk, the disk should be larger than the separation of the companions. Systems that result from turbulent fragmentation have no such expectations for their circum-multiple emission and may not have well-organized kinematics surrounding the system. We discuss our interpretation on the relative frequency of these mechanisms based on our ALMA molecular line and continuum data.

4.2.1. The Likelihood of Disk Fragmentation

The ALMA observations of close multiple protostar systems reveal several features that were previously unclear from the VLA-only data set. First, the continuum emission from extended circumbinary (or triple) material is detected in nine out of the 18 observed sources. The Class 0s and most deeply embedded Class Is have the highest tendency for a circum-multiple dust continuum. Furthermore, this circum-multiple continuum emission is almost always extended orthogonal to the outflows (L1448 IRS3B, L1448 IRS3C, L1448 IRS2, Per-emb-12, Per-emb-17, Per-emb-18), barring a few cases that appear to have inclinations that are nearly face-on (SVS13A, Per-emb-2). Second, the molecular line emission detected toward these protostars and their circum-multiple disks strongly tends to have velocity gradients orthogonal to the outflow direction (Per-emb-2, L1448 IRS3B, L1448 IRS3C, L1448 IRS2, Per-emb-17, Per-emb-18, SVS13A, L1448 IRS1). The line emission is spatially coincident with the circum-multiple dust emission (taking into account the possible dust opacity and spatial filtering in some sources), and the velocity gradients orthogonal to the outflows on these scales are interpreted as rotation because the mass of the protostar and its disk are dominant at these scales (<500 au) and do not appear to be significantly confused with large-scale effects in the clouds.

Therefore, it is clear that for many of the youngest multiple systems we have circum-multiple dust emission oriented in a manner that is consistent with a disk, and we detect molecular line kinematics that are consistent with rotation on the same scale as the circum-multiple material. We detail the molecular line detections and their likely kinematic origin for each source in Table 4. The observational results of rotating, circum-multiple gas align with the expectations for companion formation via disk fragmentation. Direct evidence for the circum-multiple material in only the younger sources makes sense, because as they evolve, the circum-multiple material is accreted rapidly (Bate 2018). Thus, about seven protostars in the sample appear to have recently undergone disk fragmentation (those listed above with velocity gradients orthogonal to the outflow and extended circum-multiple mass). Many of the Class I close systems without circum-multiple dust or gas may have formed earlier and/or consumed/dispersed their surrounding mass. The widest systems that appear convincingly consistent with disk fragmentation are Per-emb-22 (225 au) and Per-emb-33 (79, 238 au) and there is circum-multiple continuum detected in both cases. L1448 IRS1 is possibly consistent with disk fragmentation, but the companion is at a wider separation (427 au) and the disk would have had to be quite large initially to form the companion, unless there was a three-body interaction that moved the companion to a larger orbit (Reipurth & Mikkola 2012). This would then require that L1448 IRS1-A has an unresolved companion.

These ALMA observations provide important evidence for the possibility of fragmentation via gravitational instability in these systems. Additional follow-up with higher S/N will be required to confirm that these apparent disks are rotationally supported and to determine the outer radius of the Keplerian disk, ensuring that it encompasses the companion. The short observations presented here do not have the S/N to robustly fit a rotation curve that can distinguish Keplerian rotation from other velocity profiles. Therefore, at this point, we can conclude that there is evidence in favor of these systems forming via fragmentation in a gravitationally unstable disk, like L1448 IRS3B. There is rotation, circum-multiple dust, and the companions are most often orthogonal to the outflow and within the circum-multiple material.

4.2.2. Evidence for Turbulent Fragmentation in Some Systems

There are a few systems where disk fragmentation is not obviously consistent with the observations and turbulent fragmentation is more likely (i.e., Per-emb-27; Tobin et al. 2015a). Per-emb-27 in particular has some evidence for a velocity gradient in C¹⁸O but not in other tracers. H₂CO in fact appears very centrally peaked toward Per-emb-27-A and may be tracing the hot corino (e.g., Taquet et al. 2015). Furthermore, the disk toward Per-emb-27-A appears to be very small, ∼10 au in radius (Segura-Cox et al. 2018). Thus, Per-emb-27 seems unlikely to have any characteristics that make disk fragmentation a likely possibility for this system.

The source Per-emb-12 (NGC 1333 IRAS4A) is borderline for consideration as a close companion with its ∼549 au separation. We see a significant amount of material surrounding the two sources, but this is not clearly organized as a disk and is perhaps more likely a circumbinary envelope. The outflows from these two sources are roughly aligned, indicating that envelope rotation could have played a role in their formation, but the kinematics are unclear because the line emission is so complex due to the high opacity and spatial filtering that we cannot clearly say if there is organized rotation or if the inner envelope kinematics appear random.

The source Per-emb-36 (NGC 1333 IRAS2B) has extended gas emission surrounding it, but not obvious circum-multiple dust emission. Some of the gas appears off to one side and the companion is located toward the western side where much of the gas emission tends to originate. The kinematic structure of this system is unclear and the outflow is quite wide (Plunkett et al. 2013). Thus, we cannot firmly put this binary system in either category. Given that Per-emb-36 is a Class I protostar, it is possible that much of the original gas has been accreted and there is no longer a prominent circum-multiple disk. This could be a case of turbulent fragmentation, but it is hard to be more definitive with the more evolved state of Per-emb-36.

Per-emb-35 is another case that is not obviously disk fragmentation given its wide separation of 572 au. However, unlike Per-emb-12, it is much less embedded and the continuum emission is dominated by the compact circumstellar components that may exhibit rotation in some molecular species (Figure 17). While there is not much evidence for circumbinary material, the outflows (and possibly rotation axes traced by molecular lines) appear roughly parallel. Furthermore, the outflows are inclined in the same direction with respect to the plane of the sky. This case is also not obviously turbulent fragmentation or disk fragmentation. On the other hand the fragmentation mechanism in this system could be related to the overall rotation of the envelope (e.g., Burkert & Bodenheimer 1993), despite this mechanism not being favored. However, turbulent fragmentation does not rule out some outflows being parallel, just that the distribution over a larger sample should be random.

4.2.3. Which is the Most Common Mechanism?

In total, we can make reasonably strong statements about the formation mechanisms for nine out of the 17 protostars in our sample based on the extended dust emission and gas kinematics. Seven out of the nine (Per-emb-2, L1448 IRS3B, L1448 IRS3C, L1448 IRS2, Per-emb-17, Per-emb-18, SVS13A) appear to be consistent with fragmentation of a gravitationally unstable disk given their observed morphologies and velocity gradients. There are only two systems (Per-emb-12, Per-emb-27) for which we can state with reasonable confidence that they are inconsistent with disk fragmentation and that turbulent fragmentation is more likely. With the caveats of having a large separation and only a hint of rotation in ¹²CO we could add L1448 IRS1 to the disk fragmentation group. Then Per-emb-36 and Per-emb-35 could also be added to the turbulent fragmentation group, bringing the total to eight systems likely forming via disk fragmentation to four systems likely forming via turbulent fragmentation. The raw numbers indicate that for multiples on scales <600 au, ∼67% form via disk fragmentation and ∼33% form via turbulent fragmentation. If we only consider systems with <300 au separations, more in line with the sizes of disks, then there are a total of 10 systems with only two inconsistent with disk fragmentation. The smaller separation systems further indicate that the closer systems may be more frequently formed via disk fragmentation. However, these percentages obviously represent small-number statistics, prohibiting us from making a strong, general conclusion on the typical origin of close multiple systems.

A dominant contribution from turbulent fragmentation would not necessarily result in the presence of circumbinary disks within which the companions are embedded, and outflows would not tend to appear aligned orthogonal to the companion position angles and/or velocity gradients.

There is currently no clear prediction for the relative fractions from theory/simulations. Offner et al. (2010) indicate that most close companions should result from turbulent fragmentation because radiative heating of the protostellar disk should suppress fragmentation. On the other hand, the simulations from Bate (2018) include radiative feedback and do find disk fragmentation, but a study of the relative population of companions in different ranges of separations has not been conducted. However, it should be noted that Offner et al. (2010) and Bate (2012) employ different initial cloud conditions (turbulence and volume density) and different assumptions for the radiative heating. Both of these factors are likely to influence the fragmentation on cloud and disk scales in both simulations. Even in the presence of radiative heating, fragmentation can still occur for systems with total masses (star+disk) of >1 M_⊙ (Kratter et al. 2008, 2010). The Offner et al. (2010) simulations may underpredict disk fragmentation because most of the systems in that simulation were <1 M_⊙, and the different treatment of radiative heating in Bate (2012, 2018) could overpredict fragmentation in this mass range.

While we suggest that our results favor disk fragmentation as a formation mechanism for these systems, there is an inescapable caveat. Even if the disks are Keplerian, there remains the possibility that initially widely separated companions could have migrated closer together (Offner et al. 2010; Sadavoy & Stahler 2017). In the process, the original disks may have been destroyed and a new disk reformed around the companions in the same orbital plane. While the outflows could be misaligned initially (Offner et al. 2016), they could be brought into alignment and appear as a single outflow as the protostars accrete from the disk surrounding them. However, there is no evidence for past multiple outflows in different directions for the systems where we would have the possibility of determining this. But the timescales over which we could detect remnant outflows might be too short. Further study of circum-multiple material surrounding systems formed via initially turbulent fragmentation may provide clues to more uniquely distinguish between primordial disk fragmentation and turbulent fragmentation with a re-formed disk. Therefore, our main result, finding that eight out of 12 sources are consistent with disk fragmentation, could be rephrased to say that 66% is an upper limit on the frequency disk fragmentation for the production of close multiples.

4.3. Non-detections of Companions

In our ALMA follow-up of the VLA-detected close multiples, we detected companions in 12 out of 15 multiple systems that we were able to resolve. We left Per-emb-2 and Per-emb-18 out of these numbers because it is not possible to resolve their companions in these observations. All Class 0 companions are detected, and only the Class I companions have some non-detections. However, in the case of the non-detections, we do marginally detect one of the components in Per-emb-40; Per-emb-55 has a tentative detection toward the companion position. Finally, Per-emb-48 could also not be self-calibrated, reducing its likelihood to detect the companion.

The fact that most of the non-detections were more evolved Class I systems makes sense given that the VLA is detecting a combination of dust and free–free emission in most cases. The non-detected component of Per-emb-48 was Per-emb-48-B; Tychoniec et al. (2018b) detected Per-emb-48-B with a declining spectral index from 6.4 cm to 8 mm, consistent with gyro-synchrotron from a stellar corona. Per-emb-55-B, the very tentative detection, had a spectral index that was rising (0.9) but more flat than expected from dust emission alone (≥2). Assuming a spectral index of 3 (implicitly assuming a dust opacity spectral index of 1), we extrapolate the 9 mm flux density of 0.07 mJy to 1.3 mm. In order to have a flux density consistent with the observations, a flux density of less than 1 mJy, ∼95% of the 9 mm emission from Per-emb-55-B must be from free–free emission. Tychoniec et al. (2018b) found that 79% of the 9 mm emission was from free–free emission; this would also point to the spectral index of the source emission being shallower than 3. Per-emb-40-B was quite faint and also appeared to have a declining spectral index in 8 mm to 1 cm; the 4.1 cm and 6.4 cm were much fainter than expected, possibly due to source variability (e.g., Dzib et al. 2013). However, Per-emb-40-B has extended dust toward it, so it is probably real, but not detected at our level of sensitivity.

The faintness of companions toward some of these close Class I protostars is consistent with the results of Harris et al. (2012). While that study examined Class II stellar systems, they found that the systems with close separations (<300 au) had a factor of 5 lower millimeter emission compared to single systems and systems with >300 au separations. Moreover, the primary was always brighter than the secondary in the cases where both components of the system were detected at millimeter wavelengths. Thus, evolution may drive the companion(s) to be significantly fainter than the primary, and the flux density of close (<300 au) systems to be lower than more widely separated and/or single systems.

It is also worth discussing the cases of some other companion sources reported toward protostars in Perseus. A possible new companion source was reported near Per-emb-12 (NGC 1333 IRAS4A), referred to as A3 in Santangelo et al. (2015), with a 1.3 mm flux density of 187 mJy. This source A3 is not detected with the VLA in VANDAM nor with our ALMA 1.3 mm data; we only detect the two known protostars and the extended emission around them. Santangelo et al. (2015) also did not detect this source at a wavelength of 3 mm. This case of A3 toward Per-emb-12 is very similar to the case of the companions (denoted MM2 and MM3) around Per-emb-27 (NGC 1333 IRAS2A) reported by Maury et al. (2014) from PdBI observations. These companions are not detected by the VLA, CARMA, or the SMA, as detailed in Tobin et al. (2015a) or by the more sensitive ALMA data presented here. It is most likely that A3 toward Per-emb-12 and MM2 and MM3 toward Per-emb-27 are either artifacts from the deconvolution process and/or additional noise resulting from dynamic range limitations of the PdBI as was previously discussed in Tobin et al. (2015a).

In a similar vein, there have been hypotheses put forward that some clearly detected continuum sources were features produced by outflow interactions (Maury et al. 2010, 2012). However, further study over the past few years has found that the continuum features that were suggested to be outflow knots in VLA 1623 are definitively real protostars (Murillo & Lai 2013; Murillo et al. 2013; Sadavoy et al. 2018; Harris et al. 2018). The other cases where continuum sources were suggested to be outflow features (i.e., L1448-mm, NGC 1333 IRAS2A, NGC 1333 IRAS4A; Maury et al. 2010, 2014; Santangelo et al. 2015) seem to be artifacts in those data given their lack of detection here with ALMA and in other works at the same and other wavelengths (Tobin et al. 2015a, 2015b).

While it seems unlikely that outflow features are significantly contributing to the detection of false companions, we note that none of the companions without detections are found along the outflow direction. Moreover, the host protostars of the three non-detections are not driving prominent outflows as viewed from the ALMA ¹²CO data. Therefore it seems unlikely that companions are being mimicked by outflow features in our case.

We had an 80% detection rate with ALMA toward the VLA-detected close multiples in Perseus; 100% for the Class 0s and 62.5% for the Class Is. The Class I non-detections based on the VANDAM survey are likely to be real companions but were below our detection limits. For these sources in particular, it might be useful to search for their companions in the near-infrared since they are observed toward sources that are not highly embedded. Thus, we can conclude that VLA-detected companions are quite robust.

5. Conclusions

We have conducted ALMA observations toward 17 VLA-detected proto-multiple systems from the VANDAM survey that are separated by less than 600 au. We observed these systems in dust continuum at 1.3 mm and the ¹²CO, ¹³CO, C¹⁸O, SO, and H₂CO molecular lines, enabling us to examine the circumbinary environments. Our main results are as follows.

1.
Of the 15 companions that could be resolved in our observations, we detect 12 confidently, with tentative detections for two additional sources. We conclude that the VANDAM multiplicity detections from Tobin et al. (2016b) had very high reliability. Though they were not the focus of this study, we also detected all wide companions in the higher-order systems around L1448 IRS3B, Per-emb-55, and Per-emb-44.
2.
We detect circum-multiple dust emission in eight out of the nine deeply embedded Class 0 sources in the sample, but none in the remaining eight more revealed Class I systems. For the sources that are closer to edge-on, this material is oriented orthogonal to the outflow direction and the companion stars are embedded within the circum-multiple material. This is consistent with the expectations for their formation within a circumbinary/multiple disk.
3.
We detect molecular line emission associated with the circum-multiple material in 12 systems (all Class 0 systems and three out of eight Class I systems). The molecular lines with discernible velocity gradients were orthogonal to the outflow directions in eight out of 12 cases (EDJ2009-269 had too tentative a detection in ¹²CO). Velocity gradients orthogonal to the outflow on these scales (<600 au) are most likely due to rotation. Therefore, eight out of 12 systems are consistent with disk fragmentation and four out of 12 appear inconsistent with disk fragmentation. However, the overall numbers are small and may not reflect a general ratio for the origin of close companions.
4.
Most of the less embedded Class I systems in our sample do not show circumbinary dust or gas emission. These systems are likely near the end point of their evolution where the circum-multiple disk is mostly accreted or dispersed.
5.
We identified two additional strong disk candidates, L1448 IRS3A and Per-emb-8. Both show extended dust emission and have molecular line emission that is consistent with rotation. Per-emb-8 now has a better defined position angle compared to what was known previously from the VLA.

We acknowledge fruitful discussions with D. Harsono, P. Sheehan, P. Hennebelle, and N. Reynolds. We also thank the anonymous referee for comments that improved the quality of the manuscript. J.J.T. acknowledges support from NSF grant AST-1814762, the Homer L. Dodge Endowed Chair, and grant 639.041.439 from the Netherlands Organisation for Scientific Research (NWO). This paper makes use of the following ALMA data: ADS/JAO.ALMA#2013.1.00031.S. Z.-Y.L. is supported in part by NSF AST-1616636 and AST-1716259 and NASA NNX14AB38G. ALMA is a partnership of ESO (representing its member states), NSF (USA) and NINS (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 by ESO, AUI/NRAO and NAOJ. The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. This research made use of APLpy, an open-source plotting package for Python hosted at http://aplpy.github.com. This research made use of Astropy, a community-developed core Python package for Astronomy (Astropy Collaboration et al. 2018) http://www.astropy.org.

Facilities: ALMA - Atacama Large Millimeter Array, VLA - Very Large Array.

Software: Astropy (http://www.astropy.org; Greenfield et al. 2013; Astropy Collaboration et al. 2018), APLpy (http://aplpy.github.com; Robitaille & Bressert 2012), CASA (http://casa.nrao.edu; McMullin et al. 2007), The IDL Astronomy User's Library (https://idlastro.gsfc.nasa.gov/).

Appendix A: Outflow Observations

Our observations toward the Perseus multiple systems also encompassed ¹²CO, which typically traces outflowing gas toward protostars. We detected outflow signatures toward 13 out of the 17 protostars observed. These detections are shown in Figures 23–25. Of the remaining four systems that do not appear to have outflows traced by ¹²CO, rotation in the disk was found toward L1448 IRS1 (Figure 19), two traced the continuum disk without a clear velocity gradient (Per-emb-48, EDJ2009-269; Figure 20), and one was not detected (Per-emb-49). In most cases, the close multiples only have one outflow that is clearly distinguished; there are a few exceptions, however. Per-emb-12 shows two approximately parallel outflows, Per-emb-27 shows two orthogonal outflows, and Per-emb-35 shows approximately parallel outflows.

**Figure 23.** ALMA ¹²CO integrated intensity maps toward all protostars where ¹²CO appears to trace outflowing gas from the protostellar system. The integrated intensity maps are displayed as red and blue contours corresponding to the integrated intensity of line emission red- and blueshifted with respect to the system velocity. The contours are overlaid on the 1.3 mm continuum image. The line emission shows evidence for a velocity gradient consistent with rotation. The redshifted contours start at Xσ and increase in Xσ increments, and the blueshifted contours start at Xσ and increase in Xσ increments. The values for σ_red and σ_blue and velocity ranges over which the line emission was summed can be found in Table 3. The beam in the images is approximately 036 × 026.
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**Figure 24.** Same as Figure 23, but toward different sources.
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**Figure 25.** Same as Figure 23, but toward different sources.
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Appendix B: Companion Protostar Separation Distributions

The recently revised distance of ∼300 pc to the Perseus molecular cloud (Ortiz-León et al. 2018; Zucker et al. 2018) has resulted in the angular separations of the detected multiple systems corresponding to larger physical separations. We have generated a revised separation distribution using the updated distance for all the the combined sample, and the classes individually, and plot the distributions in Figure 26. We have provided the updated separations in Tables 5–7. The rest of the information contained in the tables is the same as in Tobin et al. (2016b).

**Figure 26.** Histograms of companion frequency vs. separation for multiple sources in Perseus; this is an updated version of the plot shown in Tobin et al. (2016b), but with the average distance to Perseus revised to 300 pc. The top left panel shows the distribution for all sources in the sample; the top middle and top right panels break the distribution into sources that are only comprised of Class 0 and Class I protostars, respectively. The bottom left panel shows only the multiple systems comprised of Class 0 and Class I sources, the bottom middle shows the separation distribution of all systems with a Class 0 primary source, and the bottom left panel shows the same, but with a Class I primary. The systems comprised of a Class 0 and Class I protostar are not included in the Class I plot in the bottom right. Note the apparent bi-modal distribution for the full sample and Class 0 samples and the apparent deficit of wide companions for the Class I systems. In all plots, the dashed curve is the Gaussian fit to the field star separation distribution from Raghavan et al. (2010) and the vertical dotted–dashed line corresponds to the approximate resolution limit of 20 au.
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Table 5. Class 0 Multiple Systems

Source	Separation	Separation	Flux Difference	Type
	('')	(au)	(Log [F₁/F₂])
Per-emb-2	0.080 ± 0.006	24.0 ± 1.7	0.17 ± 0.15	Class 0
Per-emb-18	0.085 ± 0.004	25.6 ± 1.2	−0.02 ± 0.12	Class 0
Per-emb-5	0.097 ± 0.006	29.1 ± 1.8	0.26 ± 0.07	Class 0
L1448NW	0.251 ± 0.004	75.3 ± 1.7	0.17 ± 0.03	Class 0
Per-emb-33	0.264 ± 0.008	79.2 ± 2.3	0.12 ± 0.05	Class 0
Per-emb-17	0.278 ± 0.014	83.3 ± 4.0	0.75 ± 0.08	Class 0
Per-emb-44	0.300 ± 0.003	90.0 ± 0.9	0.48 ± 0.02	Class 0
Per-emb-27	0.620 ± 0.003	186.0 ± 0.9	0.75 ± 0.03	Class 0
Per-emb-22	0.751 ± 0.004	225.4 ± 1.3	0.51 ± 0.05	Class 0
Per-emb-33	0.795 ± 0.004	238.4 ± 1.3	0.33 ± 0.04	Class 0
Per-emb-12	1.830 ± 0.002	548.9 ± 0.5	1.04 ± 0.02	Class 0
Per-emb-11	2.951 ± 0.008	885.4 ± 2.2	0.93 ± 0.08	Class 0
Per-emb-44+SVS13A2	5.314 ± 0.004	1594.2 ± 1.2	1.78 ± 0.03	Class 0/I
Per-emb-32	6.066 ± 0.022	1820.0 ± 6.5	−0.16 ± 0.27	Class 0/I
Per-emb-33+L1448IRS3A	7.317 ± 0.004	2195.2 ± 1.2	−0.24 ± 0.02	Class 0-Class I
Per-emb-26+Per-emb-42	8.104 ± 0.005	2431.3 ± 1.6	1.97 ± 0.03	Class 0-Class I
Per-emb-11	9.469 ± 0.025	2840.6 ± 7.6	0.76 ± 0.13	Class 0
Per-emb-8+Per-emb-55	9.557 ± 0.013	2867.2 ± 3.8	2.38 ± 0.10	Class 0-Class I
Per-emb-37+EDJ2009+235	10.556 ± 0.009	3166.7 ± 2.9	1.53 ± 0.11	Class 0-Class II
Per-emb-13+IRAS4B'	10.654 ± 0.005	3196.2 ± 1.6	0.81 ± 0.02	Class 0-Class 0
Per-emb-21+Per-emb-18	13.252 ± 0.004	3975.7 ± 1.3	−0.44 ± 0.02	Class 0-Class 0
B1-bS+Per-emb-41	13.957 ± 0.014	4187.1 ± 4.2	2.08 ± 0.09	Class 0-Class 0/I
Per-emb-44+SVS13B	14.932 ± 0.002	4479.7 ± 0.7	0.40 ± 0.01	Class 0/I-Class 0
Per-emb-16+Per-emb-28	16.063 ± 0.037	4818.9 ± 11.1	−0.05 ± 0.17	Class 0-Class 0
B1-bN+B1-bS	17.395 ± 0.009	5218.4 ± 2.6	0.31 ± 0.03	Class 0-Class 0
Per-emb-33+L1448NW	21.503 ± 0.004	6450.8 ± 1.3	−0.23 ± 0.02	Class 0-Class 0
Per-emb-18+Per-emb-49	27.474 ± 0.007	8242.3 ± 2.0	0.33 ± 0.03	Class 0-Class I
Per-emb-12+Per-emb-13	29.739 ± 0.002	8921.7 ± 0.7	1.04 ± 0.01	Class 0-Class 0
Per-emb-36+Per-emb-27	31.420 ± 0.001	9426.0 ± 0.4	0.15 ± 0.01	Class 0-Class I
Per-emb-6+Per-emb-10	31.947 ± 0.005	9584.2 ± 1.6	−0.28 ± 0.03	Class 0-Class 0
Per-emb-37+EDJ2009+233	33.704 ± 0.006	10111.3 ± 1.7	−0.62 ± 0.04	Class 0-Class II
Per-emb-44+SVS13C	34.528 ± 0.001	10358.5 ± 0.4	−0.21 ± 0.01	Class 0/I-Class 0
Per-emb-32+EDJ2009+366	36.605 ± 0.010	10981.6 ± 3.1	−1.31 ± 0.11	Class 0/I-Class II

Note. This table includes Class 0 + Class 0, Class 0 + Class I, and Class 0 + Class II multiple systems. The flux difference is calculated in the 9 mm band. The distance has been revised to 300 pc for the physical separation distance.

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Table 6. Class I Multiple Systems

Source	Separation	Separation	Flux Difference
	('')	(au)	(Log [F₁/F₂])
Per-emb-36	0.311 ± 0.005	93.4 ± 1.6	0.80 ± 0.02
Per-emb-49	0.313 ± 0.009	93.8 ± 2.6	0.46 ± 0.11
Per-emb-48	0.346 ± 0.019	103.7 ± 5.7	0.06 ± 0.17
Per-emb-40	0.391 ± 0.022	117.4 ± 6.7	0.94 ± 0.19
Per-emb-55	0.618 ± 0.009	185.3 ± 2.6	0.14 ± 0.07
EDJ2009-183	1.025 ± 0.028	307.6 ± 8.3	0.45 ± 0.13
L1448IRS1	1.424 ± 0.015	427.0 ± 4.6	1.02 ± 0.09
Per-emb-35	1.908 ± 0.003	572.3 ± 0.9	0.23 ± 0.04
EDJ2009-156	3.107 ± 0.011	932.1 ± 3.3	0.17 ± 0.13
Per-emb-58+Per-emb-65	28.878 ± 0.023	8663.3 ± 6.9	−0.44 ± 0.13

Note. This table includes only Class I + Class I multiple systems. The flux difference is calculated in the 9 mm band. The distance has been revised to 300 pc for the physical separation distance.

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Table 7. Class II Multiple Systems

Source	Separation	Separation	Flux Difference
	('')	(au)	(Log [F₁/F₂])
EDJ2009-269	0.524 ± 0.007	157.3 ± 2.1	0.12 ± 0.08
EDJ2009-156	3.107 ± 0.011	932.1 ± 3.3	0.17 ± 0.13

Note. This table includes only Class II + Class II multiple systems. The flux difference is calculated in the 9 mm band. The distance has been revised to 300 pc for the physical separation distance.

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The VLA/ALMA Nascent Disk and Multiplicity (VANDAM) Survey of Perseus Protostars. VI. Characterizing the Formation Mechanism for Close Multiple Systems

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Abstract

1. Introduction