The Gas Content and Stripping of Local Group Dwarf Galaxies

The smallest galaxies in the universe today provide a window into how the first galaxies that formed in the universe have evolved. These small galaxies can only be detected locally and have been found to have a range of properties, from objects that appear to harbor only dark matter and just thousands of stars, to gas-rich dwarf galaxies with a range of stellar populations. It remains unclear what dictates the properties of the dwarf galaxies and, in particular, what quenches their gas content and star formation. This understanding is key to making progress on how these small galaxies evolve and proceed to build up larger galaxies.

In the Local Group, the gas content of dwarf galaxies has been shown to have a clear correlation with distance from the primary spiral galaxy (e.g., Einasto et al. 1974; Blitz & Robishaw 2000; Grebel et al. 2003; Grcevich & Putman 2009, hereafter GP09; Spekkens et al. 2014). This suggests that the environment of the dwarf galaxy plays an essential role in its evolution. Beyond the Local Group, however, the situation is not as clear. The SAGA survey of Milky Way analogs finds more star-forming dwarf galaxies in the halos of most of these galaxies than for the Milky Way or M31 (Geha et al. 2017). Other work finds more consistency with the Local Group, with evidence for the gas content of dwarf galaxies increasing with isolation (Bradford et al. 2015; Stierwalt et al. 2015) and a similar distribution of dwarf galaxies around spirals in a deep optical survey (Carlsten et al. 2020).

There are numerous mechanisms that can rob a dwarf galaxy of its cold gas. In the early universe, reionization is thought to be effective at removing cold gas in the smallest dark matter halos (Tollerud & Peek 2018; Kang & Ricotti 2019; Rodriguez Wimberly et al. 2019). It is possible that this mechanism is patchy in its effectiveness, and the exact redshift is still debated, but the star formation histories of some of the oldest local dwarf galaxies are consistent with this being a quenching mechanism (Weisz et al. 2014). Another mechanism is through some type of interaction with another galaxy. There is no strong evidence that a dwarf–dwarf interaction removes the cold gas from the resulting pair (Pearson et al. 2016, 2018), but the proximity of the pair to a larger galaxy is often found to result in quenched galaxies (Geha et al. 2012; Stierwalt et al. 2015). Supernova explosions and stellar winds can also act to remove gas, particularly during periods of intense star formation, but this appears to be insufficient for the complete removal of gas in typical systems (Agertz et al. 2013; Emerick et al. 2016).

In the Local Group, dwarf galaxies within the virial radius of the Milky Way or Andromeda are largely devoid of gas, while those beyond typically have gas. This relationship suggests that the distance to a larger galaxy is an important factor in dwarf galaxy quenching. Though tidal forces can be important for those galaxies that have plunged deep into the host dark matter halo in the past, this has been shown to affect only a very small fraction of the Local Group dwarfs (e.g., Mateo et al. 2008; Simpson et al. 2018; Iorio et al. 2019). The results thus far have been consistent with a diffuse gaseous halo medium playing an important role in stripping the gas from the dwarf galaxies given the range of distances for the dwarf galaxies and the assumption of instantaneous stripping (GP09; Gatto et al. 2013). This can be tested with the discovery of new dwarf galaxies, further information about the orbital histories of the dwarf galaxies, and deeper limits on the H i content of the dwarf galaxies.

One aspect that has not been commonly considered is the role of a diffuse group medium of the Local Group in stripping dwarf galaxies. Groups are thought to harbor a hot halo medium at temperatures of ∼10^5–6 K (Osmond & Ponman 2004; Nuza et al. 2014; Stocke et al. 2019), and dwarf galaxies falling through this medium with a fast enough velocity are likely to be partially stripped. Simulations also show support for some dwarf galaxies being more easily quenched in a Local Group environment compared to the equivalent satellite of an isolated spiral (Garrison-Kimmel et al. 2019). However, the role of a large-scale environment, instead of that of an individual host galaxy, in dictating a galaxy's gas content continues to be debated (Tollerud et al. 2016; Luber et al. 2019).

In this paper, we provide new limits on the gas content of the Local Group dwarf galaxy population for galaxies discovered within 2 Mpc as of 2020 (Simon 2019, hereafter S19; McConnachie 2012, hereafter M12). This work is motivated by the discovery of large numbers of new Local Group dwarf galaxies, the availability of deeper H i survey data, and uncertainty about the cause of dwarf galaxy quenching. Gas content can be used to describe the H i results, as molecular and ionized gas are only detected in those galaxies with H i and are generally a small fraction of the H i (Gallagher et al. 2003; Buyle et al. 2006; Schruba et al. 2012). In Section 2, we describe the data and methods we use to set the H i mass limits. A model for the structure of the Local Group is presented in Section 3.1, and the Milky Way model used for orbit integration of the dwarf galaxies is described in Section 3.2. In Section 4, we present the H i results in the context of the distance to the Milky Way or Andromeda and relative to a virial radius of the Local Group. We discuss our results in Section 5 in terms of the recent Gaia proper-motion measurements for Milky Way dwarf galaxies and the potential role of a Local Group medium. The paper is briefly summarized in Section 6.

The dwarf galaxies included in this study are listed in Table 1, sorted by their heliocentric distance. This distance is also used to calculate the H i mass limit. Additional properties of the dwarf galaxies are listed in Table 2. The dwarf galaxies and their properties are primarily taken from the updated online catalog ⁶ of the M12 paper and the S19 review. The online M12 catalog contains those dwarf galaxies within 3 Mpc that have distances measured from resolved stellar populations. Here we include only those galaxies within 2 Mpc to aim for greater completeness and avoid nearby groups. This catalog includes some galaxies that are clearly not dwarf galaxies, so we exclude galaxies with stellar masses >2 × 10⁹ M_⊙. Note that this limit results in the inclusion of the Magellanic Clouds. We also exclude Canis Major from the M12 list, as it is primarily defined as a stellar stream (Martin et al. 2004). We add to the list additional ultrafaint dwarf galaxies tabulated in S19 and Antlia 2 (Torrealba et al. 2019). For all of the ultrafaint dwarf galaxies, we prioritize the more recent S19 values for their properties over M12. We examined the catalog of Karachentsev et al. (2013) as a check of our dwarf galaxy completeness within 2 Mpc and found anything additional in their catalog to be consistent with being unconfirmed or a globular cluster.

Table 1. Dwarf Galaxies within 2 Mpc: Positions and H i Parameters

Galaxy	R.A. and Decl.	V⊙	D⊙	Data a	Method b	1σ	${M}_{{\rm{H}}{\rm\small{I}}}$ c	References d
		(km s−1)	(kpc)			(mK)	(M⊙)
Draco II	15h52m48s +$64^\circ 33^{\prime} 55^{\prime\prime}$	−343	22 ± 0.4	HI4PI	Unres.	17	<57 ± 2	1, 2, 2
Segue 1	10h07m04s +$16^\circ 04^{\prime} 55^{\prime\prime}$	209	23 ± 2	GALFA1	Res.	17	<47 ± 8	1, 2, 2
DES J0225+0304	02h25m42s +$03^\circ 04^{\prime} 10^{\prime\prime}$	⋯	24 ± 1	GALFA2	Unres. no V	97	<71 ± 4	3, ⋯, 3
Tucana III	23h56m36s −$59^\circ 36^{\prime} 00^{\prime\prime}$	−102	25 ± 2	HI4PI	Unres.	20	<90 ± 14	1, 2, 2
Sagittarius dSph e	18h55m20s −$30^\circ 32^{\prime} 43^{\prime\prime}$	139	27 ± 1	HI4PI	Unres.	25	<126 ± 12	1, 2, 2
Hydrus I	02h29m33s −$79^\circ 18^{\prime} 32^{\prime\prime}$	80	28 ± 1	HI4PI	Unres.	23	<125 ± 5	4, 2, 2
Carina III	07h38m31s −$57^\circ 53^{\prime} 59^{\prime\prime}$	285	28 ± 1	HI4PI	Unres.	18	<101 ± 4	5, 2, 2
Triangulum II	02h13m17s +$36^\circ 10^{\prime} 42^{\prime\prime}$	−382	28 ± 2	HI4PI	Unres.	19	<107 ± 12	1, 2, 2
Cetus II	01h17m53s −$17^\circ 25^{\prime} 12^{\prime\prime}$	⋯	30 ± 3	HI4PI	Unres. no V	20	<127 ± 26	1, ⋯, 2
Reticulum 2	03h35m42s −$54^\circ 02^{\prime} 57^{\prime\prime}$	63	32 ± 2	HI4PI	Unres.	21	$\lt {145}_{-13}^{+14}$	1, 2, 2
Ursa Major II	08h51m30s +$63^\circ 07^{\prime} 48^{\prime\prime}$	−117	35 ± 2	HI4PI	Res.	13	$\lt {113}_{-12}^{+13}$	1, 2, 2
Carina II	07h36m26s −$57^\circ 59^{\prime} 57^{\prime\prime}$	477	36 ± 1	HI4PI	Unres.	24	<221 ± 7	5, 2, 2
Segue II	02h19m16s +$20^\circ 10^{\prime} 31^{\prime\prime}$	−40	37 ± 3	GALFA2	Res.	43	<306 ± 50	1, 2, 2
Boötes II	13h58m00s +$12^\circ 51^{\prime} 00^{\prime\prime}$	−117	42 ± 1	GALFA2	Res.	20	<185 ± 9	1, 2, 2
Coma Berenices	12h26m59s +$23^\circ 54^{\prime} 15^{\prime\prime}$	98	42 ± 2	GALFA1	Res.	12	<105 ± 8	1, 2, 2
Willman 1 f	10h49m21s +$51^\circ 03^{\prime} 00^{\prime\prime}$	−14	45 ± 10	HI4PI	Unres.	360	<5.2 ± 2.3 × 103	1, 2, 2
Pictor II	06h44m43s −$59^\circ 53^{\prime} 49^{\prime\prime}$	⋯	${45}_{-4}^{+5}$	HI4PI	Unres. no V	19	$\lt {277}_{-49}^{+62}$	6, ⋯, 2
Boötes III e	13h57m12s +$26^\circ 48^{\prime} 00^{\prime\prime}$	198	47 ± 2	GALFA1	Unres.	38	<108 ± 10	1, 1, 1
Tucana IV	00h02m55s −$60^\circ 51^{\prime} 00^{\prime\prime}$	16	48 ± 4	HI4PI	Unres.	15	<238 ± 40	7, 7, 2
LMC	05h23m34s −$69^\circ 45^{\prime} 22^{\prime\prime}$	262	51 ± 2	⋯	⋯	⋯	4.6 ± 0.4 × 108	1, 1, 1
Grus II	22h04m05s −$46^\circ 26^{\prime} 24^{\prime\prime}$	−110	53 ± 5	HI4PI	Unres.	20	<398 ± 75	7, 7, 2
Tucana V	23h37m24s −$63^\circ 16^{\prime} 12^{\prime\prime}$	−36	55 ± 9	HI4PI	Unres.	17	<360 ± 118	7, 7, 2
Tucana II	22h51m55s −$58^\circ 34^{\prime} 08^{\prime\prime}$	−129	58 ± 8	HI4PI	Res.	9	<215 ± 59	1, 2, 2
SMC	00h52m45s −$72^\circ 49^{\prime} 43^{\prime\prime}$	146	64 ± 4	⋯	⋯	⋯	4.6 ± 0.5 × 108	1, 1, 1
Boötes	14h00m06s +$14^\circ 30^{\prime} 00^{\prime\prime}$	102	66 ± 2	GALFA2	Res.	25	<564 ± 34	1, 2, 2
Sagittarius II	19h52m40s −$22^\circ 04^{\prime} 05^{\prime\prime}$	⋯	70 ± 2	HI4PI	Unres. no V	20	<696 ± 46	1, ⋯, 2
Ursa Minor	15h09m08s +$67^\circ 13^{\prime} 21^{\prime\prime}$	−247	76 ± 4	HI4PI	Unres.	15	<601 ± 63	1, 2, 2
Horologium II	03h16m32s −$50^\circ 01^{\prime} 05^{\prime\prime}$	⋯	${78}_{-7}^{+8}$	HI4PI	Unres. no V	18	$\lt {775}_{-139}^{+159}$	1, ⋯, 2
Draco	17h20m12s +$57^\circ 54^{\prime} 55^{\prime\prime}$	−291	82 ± 6	HI4PI	Res.	12	<592 ± 87	1, 2, 2
Phoenix 2	23h39m59s −$54^\circ 24^{\prime} 22^{\prime\prime}$	⋯	84 ± 4	HI4PI	Unres. no V	17	<855 ± 81	1, ⋯, 2
Sculptor g	01h00m09s −$33^\circ 42^{\prime} 33^{\prime\prime}$	111	86 ± 5	HI4PI	Unres.	62	<3.2 ± 0.4 × 103	1, 2, 2
Horologium 1	02h55m32s −$54^\circ 07^{\prime} 08^{\prime\prime}$	113	${87}_{-11}^{+13}$	HI4PI	Unres.	17	$\lt {917}_{-232}^{+274}$	1, 2, 2
Eridanus 3	02h22m46s −$52^\circ 17^{\prime} 01^{\prime\prime}$	⋯	87 ± 8	HI4PI	Unres. no V	20	<1.1 ± 0.2 × 103	1, ⋯, 1
Virgo I	12h00m10s +$00^\circ 40^{\prime} 48^{\prime\prime}$	⋯	${87}_{-8}^{+13}$	GALFA2	Unres. no V	82	$\lt {810}_{-149}^{+242}$	8, ⋯, 2
Reticulum III	03h45m26s −$60^\circ 27^{\prime} 00^{\prime\prime}$	⋯	92 ± 13	HI4PI	Unres. no V	17	<1.0 ± 0.3 × 103	1, ⋯, 2
Sextans	10h13m03s −$01^\circ 36^{\prime} 53^{\prime\prime}$	224	95 ± 3	HI4PI	Res.	8	<479 ± 30	1, 2, 2
Ursa Major I	10h34m53s +$51^\circ 55^{\prime} 12^{\prime\prime}$	−55	97 ± 6	HI4PI	Res.	103	<6.9 ± 0.8 × 103	1, 2, 2
Indus I	21h08m49s −$51^\circ 09^{\prime} 56^{\prime\prime}$	⋯	100 ± 9	HI4PI	Unres. no V	20	<1.4 ± 0.3 × 103	1, ⋯, 1
Carina	06h41m37s −$50^\circ 57^{\prime} 58^{\prime\prime}$	223	106 ± 5	HI4PI	Unres.	14	<1.1 ± 0.1 × 103	1, 2, 2
Aquarius II	22h33m56s −$09^\circ 19^{\prime} 39^{\prime\prime}$	−71	108 ± 3	HI4PI	Unres.	21	<1.8 ± 0.1 × 103	9, 2, 2
Crater II	11h49m14s −$18^\circ 24^{\prime} 47^{\prime\prime}$	88	118 ± 1	HI4PI	Res.	31	<3.0 ± 0.06 × 103	10, 2, 2
Grus I	22h56m42s −$50^\circ 09^{\prime} 48^{\prime\prime}$	−141	${120}_{-11}^{+12}$	HI4PI	Unres.	21	<2.2 ± 0.4 × 103	1, 2, 2
Pictoris 1	04h43m47s −$50^\circ 16^{\prime} 59^{\prime\prime}$	⋯	${126}_{-16}^{+19}$	HI4PI	Unres. no V	20	$\lt {2.3}_{-0.6}^{+0.7}\times {10}^{3}$	1, ⋯, 2
Antlia 2	09h35m33s −$36^\circ 46^{\prime} 02^{\prime\prime}$	291	132 ± 6	HI4PI	Res.	0.9	<111 ± 10	11, 11, 11
Hercules	16h31m02s +$12^\circ 47^{\prime} 30^{\prime\prime}$	45	132 ± 6	HI4PI	Unres.	12	<1.5 ± 0.1 × 103	1, 2, 2
Fornax h	02h39m59s −$34^\circ 26^{\prime} 57^{\prime\prime}$	55	139 ± 3	HI4PI	Unres.	25	<3.4 ± 0.1 × 103	1, 2, 2
Hydra II	12h21m42s −$31^\circ 59^{\prime} 07^{\prime\prime}$	303	${151}_{-7}^{+8}$	HI4PI	Unres.	20	<3.2 ± 0.3 × 103	1, 2, 2
Leo IV	11h32m57s +$00^\circ 32^{\prime} 00^{\prime\prime}$	132	154 ± 5	GALFA2	Res.	19	<2.3 ± 0.1 × 103	1, 2, 2
Canes Venatici II	12h57m10s +$34^\circ 19^{\prime} 15^{\prime\prime}$	−129	160 ± 4	HI4PI	Unres.	13	<2.4 ± 0.1 × 103	1, 2, 2
Leo V	11h31m10s +$02^\circ 13^{\prime} 12^{\prime\prime}$	171	169 ± 4	GALFA2	Unres.	79	<2.9 ± 0.1 × 103	1, 2, 2
Pisces II	22h58m31s +$05^\circ 57^{\prime} 09^{\prime\prime}$	−227	183 ± 15	GALFA1	Unres.	37	<1.6 ± 0.3 × 103	1, 2, 2
Columba I	05h31m26s −$28^\circ 01^{\prime} 48^{\prime\prime}$	⋯	183 ± 10	HI4PI	Unres. no V	22	<5.2 ± 0.6 × 103	1, ⋯, 2
Pegasus III	22h24m23s +$05^\circ 25^{\prime} 12^{\prime\prime}$	−223	205 ± 20	GALFA1	Unres.	57	<3.1 ± 0.6 × 103	1, 2, 2
Canes Venatici I	13h28m04s +$33^\circ 33^{\prime} 21^{\prime\prime}$	31	211 ± 6	HI4PI	Unres.	15	<4.8 ± 0.3 × 103	1, 2, 2
Indus II i	20h38m53s −$46^\circ 09^{\prime} 36^{\prime\prime}$	⋯	214 ± 16	HI4PI	Unres. no V	19	<6.2 ± 0.9 × 103	1, ⋯, 2
Leo II	11h13m29s +$22^\circ 09^{\prime} 06^{\prime\prime}$	78	233 ± 14	HI4PI	Unres.	12	<4.5 ± 0.5 × 103	1, 2, 2
Cetus III	02h05m19s −$04^\circ 16^{\prime} 12^{\prime\prime}$	⋯	${251}_{-11}^{+24}$	HI4PI	Unres. no V	19	$\lt {8.4}_{-1.0}^{+1.6}\times {10}^{3}$	12, ⋯, 2
Leo I	10h08m28s +$12^\circ 18^{\prime} 23^{\prime\prime}$	283	${254}_{-15}^{+16}$	HI4PI	Unres.	22	<1.0 ± 0.1 × 104	1, 2, 2
Eridanus II	03h44m21s −$43^\circ 32^{\prime} 00^{\prime\prime}$	76	366 ± 17	HI4PI	Unres.	19	<1.8 ± 0.2 × 104	1, 2, 2
Leo T	09h34m53s +$17^\circ 03^{\prime} 05^{\prime\prime}$	38	${409}_{-27}^{+29}$	⋯	⋯	⋯	2.8 ± 0.4 × 105	1, 2, 2
Phoenix j	01h51m06s −$44^\circ 26^{\prime} 41^{\prime\prime}$	−13	415 ± 19	⋯	⋯	⋯	1.2 ± 0.1 × 105	1, 1, 1
NGC 6822	19h44m57s −$14^\circ 47^{\prime} 21^{\prime\prime}$	−55	459 ± 17	⋯	⋯	⋯	1.3 ± 0.1 × 108	1, 1, 1
Andromeda XVI	00h59m30s +$32^\circ 22^{\prime} 36^{\prime\prime}$	−367	${476}_{-31}^{+42}$	GALFA2	Unres.	69	$\lt {2.0}_{-0.3}^{+0.4}\times {10}^{4}$	1, 1, 1
Andromeda XXIV	01h18m30s +$46^\circ 21^{\prime} 58^{\prime\prime}$	−128	600 ± 33	HI4PI	Unres.	16	<4.0 ± 0.4 × 104	1, 1, 1
NGC 185	00h38m58s +$48^\circ 20^{\prime} 15^{\prime\prime}$	−204	617 ± 26	⋯	⋯	⋯	1.1 ± 0.1 × 105	1, 1, 1
Andromeda XV	01h14m19s +$38^\circ 07^{\prime} 03^{\prime\prime}$	−323	${625}_{-35}^{+75}$	HI4PI	Unres.	15	<4.1 ± 1 × 104	1, 1, 1
Andromeda II	01h16m30s +$33^\circ 25^{\prime} 09^{\prime\prime}$	−192	652 ± 18	GALFA2	Res.	20	<4.4 ± 0.2 × 104	1, 1, 1
Andromeda XXVIII	22h32m41s +$31^\circ 12^{\prime} 58^{\prime\prime}$	−326	${661}_{-61}^{+152}$	GALFA1	Unres.	31	$\lt {1.8}_{-0.3}^{+0.8}\times {10}^{4}$	1, 1, 1
Andromeda X	01h06m34s +$44^\circ 48^{\prime} 16^{\prime\prime}$	−164	${670}_{-40}^{+25}$	HI4PI	Unres.	11	$\lt {3.4}_{-0.4}^{+0.3}\times {10}^{4}$	1, 1, 1
NGC 147	00h33m12s +$48^\circ 30^{\prime} 32^{\prime\prime}$	−193	676 ± 28	HI4PI	Unres.	21	<6.7 ± 0.6 × 104	1, 1, 1
Andromeda XXX	00h36m35s +$49^\circ 38^{\prime} 48^{\prime\prime}$	−140	${682}_{-82}^{+31}$	HI4PI	Unres.	12	$\lt {4.0}_{-1.0}^{+0.4}\times {10}^{4}$	1, 1, 1
Andromeda XVII	00h37m07s +$44^\circ 19^{\prime} 20^{\prime\prime}$	−252	${728}_{-27}^{+37}$	HI4PI	Unres.	11	$\lt {4.2}_{-0.3}^{+0.4}\times {10}^{4}$	1, 1, 1
Andromeda XXIX	23h58m56s +$30^\circ 45^{\prime} 20^{\prime\prime}$	−194	731 ± 74	GALFA1	Unres.	36	<2.5 ± 0.5 × 104	1, 1, 1
Andromeda XI	00h46m20s +$33^\circ 48^{\prime} 05^{\prime\prime}$	−420	735 ± 17	HI4PI	Unres.	12	<4.6 ± 0.2 × 104	1, 1, 1
Andromeda XX	00h07m31s +$35^\circ 07^{\prime} 56^{\prime\prime}$	−456	${741}_{-55}^{+41}$	GALFA2	Unres.	76	$\lt {5.4}_{-0.8}^{+0.6}\times {10}^{4}$	1, 1, 1
Andromeda I	00h45m40s +$38^\circ 02^{\prime} 28^{\prime\prime}$	−376	745 ± 24	HI4PI	Unres.	19	<7.5 ± 0.5 × 104	1, 1, 1
Andromeda III	00h35m34s +$36^\circ 29^{\prime} 52^{\prime\prime}$	−344	748 ± 24	GALFA2	Unres.	86	<6.2 ± 0.4 × 104	1, 1, 1
IC 1613	01h04m48s +$02^\circ 07^{\prime} 04^{\prime\prime}$	−232	755 ± 42	⋯	⋯	⋯	6.5 ± 0.7 × 107	1, 1, 1
Cetus	00h26m11s −$11^\circ 02^{\prime} 40^{\prime\prime}$	−84	755 ± 24	HI4PI	Unres.	24	<9.5 ± 0.6 × 104	1, 1, 1
Lacerta I k	22h58m16s +$41^\circ 17^{\prime} 28^{\prime\prime}$	−198	759 ± 42	HI4PI	Unres.	24	<9.6 ± 1.1 × 104	1, 13, 1
Andromeda VII	23h26m32s +$50^\circ 40^{\prime} 33^{\prime\prime}$	−307	762 ± 35	HI4PI	Unres.	16	<6.6 ± 0.6 × 104	1, 1, 1
Andromeda XXVI	00h23m46s +$47^\circ 54^{\prime} 58^{\prime\prime}$	−262	762 ± 42	HI4PI	Unres.	12	<5.0 ± 0.6 × 104	1, 1, 1
Andromeda IX	00h52m53s +$43^\circ 11^{\prime} 45^{\prime\prime}$	−209	766 ± 25	HI4PI	Unres.	13	<5.5 ± 0.4 × 104	1, 1, 1
LGS 3	01h03m55s +$21^\circ 53^{\prime} 06^{\prime\prime}$	−287	769 ± 25	⋯	⋯	⋯	3.8 ± 0.2 × 105	1, 1, 1
Andromeda XXIII	01h29m22s +$38^\circ 43^{\prime} 08^{\prime\prime}$	−238	769 ± 46	HI4PI	Unres.	20	<8.2 ± 1 × 104	1, 1, 1
Perseus I l	03h01m24s +$40^\circ 59^{\prime} 18^{\prime\prime}$	−326	773 ± 64	HI4PI	Unres.	19	<7.9 ± 1.3 × 104	1, 13, 1
Andromeda V	01h10m17s +$47^\circ 37^{\prime} 41^{\prime\prime}$	−403	773 ± 29	HI4PI	Unres.	19	<8.0 ± 0.6 × 104	1, 1, 1
Cassiopeia III m	00h35m59s +$51^\circ 33^{\prime} 35^{\prime\prime}$	−372	776 ± 50	HI4PI	Unres.	18	<7.6 ± 1.0 × 104	1, 13, 1
Andromeda VI	23h51m46s +$24^\circ 34^{\prime} 57^{\prime\prime}$	−340	783 ± 25	GALFA1	Unres.	30	<2.4 ± 0.2 × 104	1, 1, 1
Andromeda XIV	00h51m35s +$29^\circ 41^{\prime} 49^{\prime\prime}$	−481	${794}_{-205}^{+22}$	GALFA1	Unres.	29	$\lt {2.4}_{-1.2}^{+0.1}\times {10}^{4}$	1, 1, 1
IC 10	00h20m17s +$59^\circ 18^{\prime} 14^{\prime\prime}$	−348	794 ± 44	⋯	⋯	⋯	5.0 ± 0.6 × 107	1, 1, 1
Leo A	09h59m26s +$30^\circ 44^{\prime} 47^{\prime\prime}$	24	798 ± 44	⋯	⋯	⋯	1.1 ± 0.1 × 107	1, 1, 1
M32	00h42m42s +$40^\circ 51^{\prime} 55^{\prime\prime}$	−199	805 ± 78	HI4PI	Unres.	38	<1.8 ± 0.3 × 105	1, 1, 1
Andromeda XXV	00h30m09s +$46^\circ 51^{\prime} 07^{\prime\prime}$	−108	813 ± 45	HI4PI	Unres.	62	<2.9 ± 0.3 × 105	1, 1, 1
Andromeda XIX	00h19m32s +$35^\circ 02^{\prime} 37^{\prime\prime}$	−112	${820}_{-162}^{+30}$	GALFA2	Res.	39	$\lt {1.3}_{-0.5}^{+0.1}\times {10}^{5}$	1, 1, 1
NGC 205	00h40m22s +$41^\circ 41^{\prime} 07^{\prime\prime}$	−246	824 ± 27	⋯	⋯	⋯	4.0 ± 0.3 × 105	1, 1, 1
Andromeda XXI	23h54m48s +$42^\circ 28^{\prime} 15^{\prime\prime}$	−363	${828}_{-27}^{+23}$	HI4PI	Unres.	14	<6.9 ± 0.4 × 104	1, 1, 1
Andromeda XXVII	00h37m27s +$45^\circ 23^{\prime} 13^{\prime\prime}$	−540	828 ± 46	HI4PI	Unres.	19	<9.1 ± 1 × 104	1, 1, 1
Andromeda XIII	00h51m51s +$33^\circ 00^{\prime} 16^{\prime\prime}$	−185	840 ± 19	GALFA2	Unres.	76	<7.0 ± 0.3 × 104	1, 1, 1
Tucana	22h41m50s −$64^\circ 25^{\prime} 10^{\prime\prime}$	194	887 ± 49	HI4PI	Unres.	15	<8.5 ± 0.9 × 104	1, 1, 1
Andromeda XXII	01h27m40s +$28^\circ 05^{\prime} 25^{\prime\prime}$	−130	${920}_{-153}^{+30}$	GALFA1	Unres.	36	$\lt {3.9}_{-1.3}^{+0.3}\times {10}^{4}$	1, 1, 1
Pegasus dIrr	23h28m36s +$14^\circ 44^{\prime} 35^{\prime\prime}$	−180	920 ± 30	⋯	⋯	⋯	5.9 ± 0.4 × 106	1, 1, 1
Andromeda XII	00h47m27s +$34^\circ 22^{\prime} 29^{\prime\prime}$	−558	${929}_{-145}^{+39}$	GALFA2	Unres.	71	$\lt {7.9}_{-2.5}^{+0.7}\times {10}^{4}$	1, 1, 1
WLM	00h01m58s −$15^\circ 27^{\prime} 39^{\prime\prime}$	−130	933 ± 34	⋯	⋯	⋯	6.1 ± 0.5 × 107	1, 1, 1
Sagittarius dIrr	19h29m59s −$17^\circ 40^{\prime} 51^{\prime\prime}$	−79	1067 ± 88	⋯	⋯	⋯	8.8 ± 1.5 × 106	1, 1, 1
Aquarius	20h46m52s −$12^\circ 50^{\prime} 53^{\prime\prime}$	−138	1072 ± 40	⋯	⋯	⋯	4.1 ± 0.3 × 106	1, 1, 1
Andromeda XVIII	00h02m14s +$45^\circ 05^{\prime} 20^{\prime\prime}$	−332	${1213}_{-45}^{+39}$	HI4PI	Unres.	21	$\lt {2.2}_{-0.2}^{+0.1}\times {10}^{5}$	1, 1, 1
Antlia B	09h48m56s −$25^\circ 59^{\prime} 24^{\prime\prime}$	376	1294 ± 95	⋯	⋯	⋯	2.8 ± 0.4 × 105	1, 1, 1
NGC 3109	10h03m07s −$26^\circ 09^{\prime} 35^{\prime\prime}$	403	1300 ± 48	⋯	⋯	⋯	4.5 ± 0.3 × 108	1, 1, 1
Antlia	10h04m04s −$27^\circ 19^{\prime} 52^{\prime\prime}$	362	1349 ± 62	⋯	⋯	⋯	7.3 ± 0.7 × 105	1, 1, 1
UGC 4879	09h16m02s +$52^\circ 50^{\prime} 24^{\prime\prime}$	−29	1361 ± 25	⋯	⋯	⋯	9.5 ± 0.4 × 105	1, 1, 1
Sextans B	10h00m00s +$05^\circ 19^{\prime} 56^{\prime\prime}$	304	1426 ± 20	⋯	⋯	⋯	5.1 ± 0.1 × 107	1, 1, 1
Sextans A	10h11m01s −$04^\circ 41^{\prime} 34^{\prime\prime}$	324	1432 ± 53	⋯	⋯	⋯	7.7 ± 0.6 × 107	1, 1, 1
Leo P	10h21m45s +$18^\circ 05^{\prime} 17^{\prime\prime}$	264	1622 ± 149	⋯	⋯	⋯	9.4 ± 1.7 × 105	1, 1, 1
HIZSS 3B	07h00m29s −$04^\circ 12^{\prime} 30^{\prime\prime}$	323	1675 ± 108	⋯	⋯	⋯	2.6 ± 0.3 × 106	1, 1, 1
HIZSS 3(A)	07h00m29s −$04^\circ 12^{\prime} 30^{\prime\prime}$	288	1675 ± 108	⋯	⋯	⋯	1.4 ± 0.2 × 107	1, 1, 1
KKR 25	16h13m48s +$54^\circ 22^{\prime} 16^{\prime\prime}$	−65	1923 ± 62	HI4PI	Unres.	23	<6.1 ± 0.4 × 105	1, 1, 1
ESO 410–G005	00h15m32s −$32^\circ 10^{\prime} 48^{\prime\prime}$	⋯	1923 ± 35	⋯	⋯	⋯	7.3 ± 0.3 × 105	1, 14, 1
IC 5152	22h02m42s −$51^\circ 17^{\prime} 47^{\prime\prime}$	122	1950 ± 45	⋯	⋯	⋯	8.7 ± 0.4 × 107	1, 1, 1

Notes.

^a The data source is either HI4PI (HI4PI Collaboration et al. 2016), GALFA-H i DR1 (GALFA1; Peek et al. 2011), or GALFA-H i DR2 (GALFA2; Peek et al. 2018). ^b The method indicates whether the 1σ was obtained with a method consistent with the source being resolved (Res.) or unresolved (Unres.) in the data. If the source does not have a velocity (no V), then the 1σ is an average from all of the channels without significant Galactic emission. For HI4PI resolved, the 1σ values tend to be lower, as they are the standard deviation of the flux sum divided by the area of the optical galaxy for nine regions surrounding the galaxy (see text). ^c The H i mass is calculated using 5σ, the noted distance, and a fixed velocity width of 10 km s⁻¹. The errors are from the distance errors shown. ^d This column notes the references for the positions, velocities, and distances, respectively. ^e The limit is only for the core of this tidally extended galaxy. ^f The limit is calculated with twice the median flux in the region of the galaxy due to strong Galactic contamination (see text). ^g Sculptor has been claimed as a detection in previous publications, but the H i clouds are offset from the stellar position of the galaxy, and there is abundant surrounding emission at similar velocities. ^h Fornax is at velocities that confuse any associated H i with Galactic emission. The limit at the stellar position and velocity is given. ⁱ This galaxy is unlikely to exist based on observations by Cantu et al. (2020). ^j Phoenix has confusion with Galactic emission, but the H i cloud is at the position and velocity of the stellar component. ^k Lacerta I is Andromeda XXXI in (M12). ^l Perseus I is Andromeda XXXIII in (M12). ^m Cassiopeia III is Andromeda XXXII in (M12).

References. (1) M12. (2) S19. (3) Luque et al. (2017). (4) Koposov et al. (2018). (5) Torrealba et al. (2017). (6) Drlica-Wagner et al. (2016). (7) Simons et al. (2020). (8) Homma et al. (2016). (9) Torrealba et al. (2016b). (10) Torrealba et al. (2016a). (11) Torrealba et al. (2019). (12) Homma et al. (2018). (13) Martin et al. (2014). (14) Bouchard et al. (2005) found a velocity of 159 km s⁻¹, but we do not note that here, as this could not be confirmed by Westmeier et al. (2017) or Koribalski et al. (2018). Koribalski et al. (2018) found emission with a similar flux and velocity width at 36 km s⁻¹, so we leave the H i mass.

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Table 2. Dwarf Galaxies within 2 Mpc: Other Parameters ^a

Galaxy	LV (L⊙)	rh (pc)	σ* (km s−1) b	DMW (kpc)	DM31 (kpc)	DLG (kpc) c	References d
Draco II	1.8 × 102	19	2.9	24	771	$-{210}_{-15}^{+15}$	2, 2, 3
Segue 1	2.8 × 102	24	3.7	28	788	$-{197}_{-17}^{+18}$	2, 2, 2
DES J0225+0304	2.4 × 102	19	⋯	29	762	$-{214}_{-15}^{+14}$	2, 2, ⋯
Tucana III	3.4 × 102	37	<1.2	23	784	$-{203}_{-17}^{+18}$	2, 2, 2
Sagittarius dSph	2.1 × 107	2662	11.4	19	787	$-{206}_{-17}^{+19}$	2, 2, 2
Hydrus I	6.5 × 103	53	2.7	26	794	$-{198}_{-17}^{+20}$	2, 2, 2
Carina III	7.8 × 102	30	5.6	29	797	$-{195}_{-18}^{+20}$	2, 2, 2
Triangulum II	3.7 × 102	16	<3.4	35	752	$-{222}_{-14}^{+12}$	2, 2, 2
Cetus II	8.6 × 101	17	⋯	32	764	$-{208}_{-14}^{+13}$	2, 2, ⋯
Reticulum 2	3.0 × 103	51	3.3	33	786	$-{193}_{-16}^{+17}$	2, 2, 2
Ursa Major II	5.1 × 103	139	5.6	41	766	$-{198}_{-14}^{+13}$	2, 2, 2
Carina II	5.4 × 103	92	3.4	37	803	$-{187}_{-18}^{+20}$	2, 2, 2
Segue II	5.3 × 102	40	3.4	43	747	$-{217}_{-13}^{+11}$	2, 2, 2
Coma Berenices	4.4 × 103	69	4.6	43	797	$-{181}_{-17}^{+18}$	2, 2, 2
Boötes II	1.3 × 103	39	10.5	40	803	$-{184}_{-18}^{+20}$	2, 2, 2
Pictor II	1.6 × 103	47	⋯	46	806	$-{179}_{-18}^{+20}$	4, 4, ⋯
Willman 1	1.2 × 103	33	4.0	50	776	$-{183}_{-15}^{+13}$	2, 2, 2
Boötes III f	1.7 × 104	⋯	10.7	46	796	$-{179}_{-17}^{+17}$	1, ⋯, 5
Tucana IV	2.1 × 103	127	4.3	45	791	$-{180}_{-16}^{+16}$	2, 2, 2
LMC	1.5 × 109	2675	σgas = 34	50	807	$-{174}_{-18}^{+19}$	1, 7, e
Grus II	3.1 × 103	93	<1.9	48	785	$-{180}_{-16}^{+15}$	2, 2, 6
Tucana V	3.7 × 102	16	⋯	52	795	$-{173}_{-16}^{+16}$	2, 2, ⋯
Tucana II	3.1 × 103	121	8.6	54	794	$-{172}_{-16}^{+16}$	2, 2, 2
SMC	4.6 × 108	1105	27.6	61	807	$-{163}_{-17}^{+18}$	1, 7, 1
Boötes	2.2 × 104	191	4.6	64	815	$-{161}_{-18}^{+20}$	2, 2, 2
Sagittarius II	1.0 × 104	33	⋯	63	785	$-{166}_{-15}^{+13}$	2, 2, ⋯
Ursa Minor	3.5 × 105	405	9.5	78	754	$-{171}_{-12}^{+8}$	2, 2, 2
Horologium II	3.6 × 102	44	⋯	79	793	$-{150}_{-15}^{+13}$	2, 2, ⋯
Draco	3.0 × 105	231	9.1	82	748	$-{171}_{-12}^{+7}$	2, 2, 2
Phoenix 2	1.0 × 103	37	⋯	81	793	$-{147}_{-15}^{+13}$	2, 2, ⋯
Sculptor	1.8 × 106	279	9.2	86	761	$-{158}_{-12}^{+8}$	2, 2, 2
Horologium 1	2.7 × 103	40	4.9	87	798	$-{141}_{-15}^{+13}$	2, 2, 2
Virgo I	1.8 × 102	38	⋯	87	845	$-{141}_{-21}^{+24}$	8, 8, ⋯
Eridanus 3	5.4 × 102	14	⋯	87	793	$-{142}_{-15}^{+12}$	1, 1, ⋯
Reticulum III	1.8 × 103	64	⋯	92	813	$-{133}_{-16}^{+15}$	2, 2, ⋯
Sextans	3.2 × 105	456	7.9	98	841	$-{127}_{-19}^{+21}$	2, 2, 2
Ursa Major	9.6 × 103	295	7.0	102	773	$-{137}_{-13}^{+8}$	2, 2, 2
Indus I	2.1 × 103	37	⋯	94	808	$-{132}_{-16}^{+14}$	1, 1, ⋯
Carina	5.2 × 105	311	6.6	108	838	$-{117}_{-18}^{+19}$	2, 2, 2
Aquarius II	4.7 × 103	160	5.4	105	729	$-{164}_{-10}^{+3}$	2, 2, 2
Crater II	1.6 × 105	1066	2.7	116	886	$-{126}_{-26}^{+35}$	2, 2, 2
Grus I	2.1 × 103	28	2.9	116	797	$-{115}_{-14}^{+10}$	2, 2, 2
Pictoris 1	2.5 × 103	32	⋯	128	822	$-{99}_{-16}^{+13}$	2, 2, ⋯
Antlia 2	3.5 × 105	2899	5.7	133	889	$-{101}_{-25}^{+30}$	9, 9, 9
Hercules	1.8 × 104	216	5.1	126	822	$-{100}_{-16}^{+13}$	2, 2, 2
Fornax	1.9 × 107	792	11.7	141	768	$-{105}_{-11}^{+5}$	2, 2, 2
Hydra II	7.5 × 103	67	<3.6	148	928	$-{118}_{-36}^{+52}$	2, 2, 2
Leo IV	8.5 × 103	114	3.3	155	895	$-{75}_{-23}^{+26}$	2, 2, 2
Canes Venatici II	1.0 × 104	71	4.6	161	833	$-{66}_{-15}^{+12}$	2, 2, 2
Leo V	4.4 × 103	49	2.3	170	904	$-{60}_{-23}^{+26}$	2, 2, 2
Pisces II	4.2 × 103	60	5.4	182	655	$-{154}_{-6}^{-4}$	2, 2, 2
Columba I	4.1 × 103	117	⋯	188	819	$-{45}_{-13}^{+7}$	2, 2, ⋯
Pegasus III	3.7 × 103	78	5.4	203	657	$-{128}_{-7}^{-4}$	2, 2, 2
Canes Venatici I	2.7 × 105	437	7.6	211	856	$-{16}_{-15}^{+10}$	2, 2, 2
Indus II	4.5 × 103	181	⋯	208	853	$-{20}_{-15}^{+10}$	2, 2, ⋯
Leo II	6.7 × 105	171	7.4	236	897	${11}_{-17}^{+14}$	2, 2, 2
Cetus III	8.2 × 102	90	⋯	255	644	$-{89}_{-10}^{-2}$	10, 10, ⋯
Leo I	4.4 × 106	270	9.2	258	918	${34}_{-18}^{+15}$	2, 2, 2
Eridanus II	5.9 × 104	246	6.9	368	884	${128}_{-10}^{+1}$	2, 2, 2
Leo T	1.4 × 105	118	7.5	414	982	${187}_{-15}^{+8}$	2, 2, 2
Phoenix	7.7 × 105	454	9.3	415	864	${160}_{-4}^{-7}$	1, 1, 11
NGC 6822	1.0 × 108	354	23.2	452	894	${200}_{-4}^{-7}$	1, 1, 1
Andromeda XVI	3.4 × 105	123	3.8	480	319	$-{231}_{-29}^{+12}$	1, 1, 1
Andromeda XXIV	9.3 × 104	366	⋯	605	204	$-{219}_{-34}^{+17}$	1, 1, ⋯
NGC 185	6.8 × 107	457	24.0	621	184	$-{228}_{-35}^{+18}$	1, 1, 1
Andromeda XV	4.8 × 105	220	4.0	630	175	$-{229}_{-35}^{+18}$	1, 1, 1
Andromeda II	9.1 × 106	1175	7.8	656	181	$-{185}_{-36}^{+19}$	1, 1, 1
Andromeda XXVIII	2.1 × 105	213	6.6	661	365	${0}_{-34}^{+16}$	1, 1, 1
Andromeda X	8.8 × 104	253	6.4	674	130	$-{233}_{-37}^{+20}$	1, 1, 1
NGC 147	6.2 × 107	623	16.0	680	139	$-{212}_{-37}^{+20}$	1, 1, 1
Andromeda XXX	1.3 × 105	268	11.8	686	144	$-{199}_{-37}^{+20}$	1, 1, 1
Andromeda XVII	2.2 × 105	263	2.9	732	66	$-{251}_{-40}^{+23}$	1, 1, 1
Andromeda XXIX	1.8 × 105	362	5.7	734	187	$-{123}_{-38}^{+21}$	1, 1, 1
Andromeda XI	4.6 × 104	152	<4.6	738	108	$-{196}_{-39}^{+23}$	1, 1, 1
Andromeda XX	2.5 × 104	114	7.1	744	128	$-{173}_{-39}^{+23}$	1, 1, 1
Andromeda I	4.7 × 106	672	10.2	749	55	$-{247}_{-40}^{+24}$	1, 1, 1
Andromeda III	1.0 × 106	479	9.3	752	73	$-{224}_{-40}^{+24}$	1, 1, 1
IC 1613	1.0 × 108	1496	10.8	758	518	${172}_{-33}^{+15}$	1, 1, 1
Cetus	2.8 × 106	703	8.3	756	678	${288}_{-29}^{+11}$	1, 1, 1
Lacerta I	4.1 × 106	927	10.3	760	262	$-{42}_{-38}^{+21}$	1, 1, 12
Andromeda XXVI	6.0 × 104	222	8.6	766	102	$-{189}_{-40}^{+24}$	1, 1, 1
Andromeda VII	1.6 × 107	776	13.0	765	217	$-{81}_{-39}^{+22}$	1, 1, 1
Andromeda IX	1.5 × 105	557	10.9	770	39	$-{249}_{-41}^{+25}$	1, 1, 1
Andromeda XXIII	1.1 × 106	1029	7.1	774	126	$-{163}_{-40}^{+24}$	1, 1, 1
LGS 3	9.6 × 105	470	7.9	773	268	$-{33}_{-39}^{+21}$	1, 1, 1
Perseus I	1.2 × 106	382	4.2	779	348	${41}_{-37}^{+20}$	1, 1, 12
Andromeda V	5.6 × 105	315	11.5	777	109	$-{178}_{-41}^{+25}$	1, 1, 1
Cassiopeia III	6.8 × 106	1468	8.4	780	140	$-{147}_{-40}^{+24}$	1, 1, 12
Andromeda VI	3.3 × 106	524	12.4	785	268	$-{28}_{-39}^{+22}$	1, 1, 1
IC 10	8.6 × 107	612	σgas = 25.9	798	252	$-{40}_{-40}^{+23}$	1, 1, 13
Andromeda XIV	2.4 × 105	393	5.3	798	161	$-{124}_{-41}^{+25}$	1, 1, 1
Leo A	6.0 × 106	499	6.7	803	1197	${564}_{-3}^{-9}$	1, 1, 1
M32	3.2 × 108	110	92.0	809	27	$-{271}_{-43}^{+28}$	1, 1, 1
Andromeda XXV	6.8 × 105	709	3.0	817	90	$-{192}_{-42}^{+27}$	1, 1, 1
Andromeda XIX	3.3 × 105	3065	7.8	824	115	$-{167}_{-42}^{+27}$	1, 14, 14
NGC 205	3.3 × 108	590	35.0	828	46	$-{263}_{-43}^{+29}$	1, 1, 1
Andromeda XXVII	1.2 × 105	433	14.8	832	77	$-{212}_{-43}^{+28}$	1, 1, 1
Andromeda XXI	7.0 × 105	843	4.5	831	135	$-{147}_{-42}^{+27}$	1, 1, 1
Andromeda XIII	3.5 × 104	191	5.8	843	134	$-{150}_{-43}^{+28}$	1, 1, 1
Tucana	5.6 × 105	284	6.2	883	1352	${657}_{-15}^{+6}$	1, 1, 15
Andromeda XXII	4.6 × 104	252	2.8	925	276	$-{11}_{-44}^{+29}$	1, 1, 1
Pegasus dIrr	6.6 × 106	562	12.3	921	474	${187}_{-40}^{+22}$	1, 1, 1
Andromeda XII	3.5 × 104	324	2.6	933	182	$-{137}_{-46}^{+34}$	1, 1, 1
WLM	4.3 × 107	2340	17.5	933	835	${476}_{-31}^{+12}$	1, 1, 1
Sagittarius dIrr	3.5 × 106	282	σgas = 10	1059	1354	${806}_{-9}^{-5}$	1, 1, 1
Aquarius	1.6 × 106	458	7.9	1066	1170	${740}_{-22}^{+5}$	1, 1, 1
Andromeda XVIII	5.0 × 105	325	9.7	1217	457	${44}_{-58}^{+56}$	1, 1, 1
Antlia B	6.3 × 105	271	σgas = 7.2	1296	1963	${1047}_{-45}^{+43}$	1, 1, 16
NGC 3109	7.6 × 107	1626	σgas = 61	1301	1984	${1049}_{-47}^{+46}$	1, 1, 17
Antlia	1.3 × 106	471	σgas = 12.7	1350	2036	${1096}_{-48}^{+47}$	1, 1, 17
UGC 4879	8.3 × 106	162	9.6	1367	1394	${1017}_{-26}^{+8}$	1, 1, 1
Sextans B	5.2 × 107	440	σgas = 17	1429	1940	${1204}_{-25}^{+18}$	1, 1, 18
Sextans A	4.4 × 107	1029	σgas = 19	1435	2024	${1201}_{-35}^{+30}$	1, 1, 18
Leo P	3.9 × 105	566	σgas = 8.4	1625	2048	${1401}_{-18}^{+8}$	1, 19, 19
HIZSS 3B	⋯	⋯	σgas = 11.9	1681	1921	${1429}_{-11}^{-4}$	⋯, ⋯, 20
HIZSS 3(A)	⋯	⋯	σgas = 23	1681	1921	${1429}_{-11}^{-4}$	⋯, ⋯, 20
KKR 25	2.0 × 106	263	⋯	1922	1869	${1541}_{-31}^{+12}$	1, 1, ⋯
ESO 410–G005	3.5 × 106	280	σgas = 14.2	1922	1861	${1536}_{-32}^{+12}$	1, 1, 21
IC 5152	1.4 × 108	550	σgas = 36	1945	2209	${1702}_{-8}^{-6}$	1, 1, 22

Notes.

^a The parameters use the distances and coordinates from Table 1 where relevant. ^b When σ_gas is noted, it is derived directly from the FWHM of the H i line. ^c This is the distance from the Local Group surface defined in Section 3.1, while D_MW and D_M31 are from the center of the Milky Way and M31, respectively. The superscript value is the variation on the distance from this surface when the Milky Way mass is higher and the M31 mass is lower. The subscript value is when the Milky Way mass is lower and the M31 mass is higher. The total Local Group mass remains the same in all cases. ^d This column notes the references for the luminosity, half-light radius, and velocity dispersion, respectively. ^e The LMC has a very large stellar rotation component compared to the stellar dispersion component. Using the stellar dispersion value (20.2 km s⁻¹) would lead to a severe underestimate of the total mass, so we use the gas dispersion calculated from the FWHM of the H i line. ^f No r_h is given, since this galaxy is highly disrupted.

References. (1) M12. (2) S19. (3) Martin et al. (2016b). (4) Drlica-Wagner et al. (2016). (5) Carlin & Sand (2018). (6) Simons et al. (2020). (7) Bothun & Thompson (1988). (8) Homma et al. (2016). (9) Torrealba et al. (2019). (10) Homma et al. (2018). (11) Kacharov et al. (2017). (12) Martin et al. (2014). (13) Namumba et al. (2019). (14) Martin et al. (2016a). Collins et al. (2020). (15) Taibi et al. (2020). (16) Sand et al. (2015). (17) Barnes et al. (2001). (18) Namumba et al. (2018). (19) Bernstein-Cooper et al. (2014). The r_h in this case is the semimajor axis noted by McQuinn et al. (2015). (20) Begum et al. (2005). (21) The σ is from Bouchard et al. (2005) and is consistent with the detection at the lower velocity by Koribalski et al. (2018). (22) Koribalski et al. (2018).

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There are three main sources of H i data we used to set limits on the dwarf galaxy H i content: Galactic Arecibo L-band Feed Array HI (GALFA-H i DR1; Peek et al. 2011), GALFA-H i DR2 ⁷ (Peek et al. 2018), and HI4PI (HI4PI Collaboration et al. 2016). These are referred to as GALFA1, GALFA2, and HI4PI in Table 1. The starting point for all of the GALFA-H i data were cubes with a resolution of 4' and a channel spacing of 0.754 km s⁻¹. GALFA-H i DR2 covers the entire sky observable with the Arecibo 300 m radio telescope, or all right ascensions for declinations between approximately −1° and +37° (approximately one-third of the sky). GALFA-H i DR1 has the same resolution as DR2 but only covers approximately half of the Arecibo sky in noncontinuous regions. We used the DR1 data in a few cases because they are deeper in some regions with the inclusion of additional commensal data (see Peek et al. 2011, 2018). The HI4PI survey combines the southern GASS data taken with the Parkes 64 m radio telescope (McClure-Griffiths et al. 2009) and the northern EBHIS survey taken with the Effelsberg 100 m telescope to cover the entire sky (Winkel et al. 2016). The combination results in an all-sky H i survey with 162 spatial resolution and a channel spacing of 1.29 km s⁻¹.

To determine if a dwarf galaxy was detected and obtain the limits on the H i masses, we first applied Hanning smoothing on each cube that contains the position of a dwarf galaxy to ∼10 km s⁻¹. Since the channel spacing is different for GALFA-H i and HI4PI, the closest value to 10 km s⁻¹ was used; 13 channels were smoothed for GALFA-H i and 7 for HI4PI. We then examined the smoothed data cube at the position and velocity of the source in both the data cube and the spectrum for a possible detection. In all cases where there was no previous detection, there was no detection at the position and velocity of the galaxy, and we proceeded to extract a noise level to determine the H i mass limit.

The method to obtain the quantitative H i mass limit differed depending on whether the gas in the dwarf was likely to be resolved or not. We used the optical half-light radius (r_h ) as a guide to whether this was likely. In all cases, the noise values were checked to be consistent across several channels, and, with the exception of the rare cases when the dwarf was resolved with HI4PI data, the standard deviation was calculated within a 30' diameter region at the position and velocity of the dwarf. The 30' window is designed to be an area larger than the beam size but not large enough to potentially encompass nearby structures at the velocity of the dwarf (i.e., Galactic emission in some cases or high-velocity clouds). None of the dwarf galaxies were close enough to the edges of the cubes for this to influence the analysis. If the dwarf was in the GALFA-H i sky and the optical diameter was smaller than 4'–5', the limit was obtained from the standard deviation within a 30' diameter region at the position and velocity of the dwarf (Unres. in Table 1 with data = GALFA1 or 2). If the dwarf was in the GALFA-H i sky and had an optical diameter between approximately 5' and 12', we did this same procedure with GALFA-H i data smoothed to 8' (Res. in Table 1 with data = GALFA1 or 2). This generally resulted in a better limit on the H i mass than going to the HI4PI data, though for the three dwarf galaxies with optical diameters >9', it does assume that the H i is more centrally concentrated (Coma Berenices, Andromeda XIX, and Andromeda II). There are several dwarf galaxies for which we adopt the HI4PI limit even though the dwarf is in the GALFA-H i sky: Hercules and Canes Venatici I because of their large optical size and Canes Venatici II, Andromeda XI, Leo I, and Leo II due to their location in a noisier region of the GALFA-H i data (i.e., on a basket-weave remnant and/or near the outer regions of the survey).

The HI4PI data were used for the majority of the dwarf galaxies due to their greater sky coverage. The limits were obtained with a 30' region (Unres. in Table 1 with data = HI4PI), as was used for the GALFA data, with the exception of those galaxies that have optical sizes larger than the 162 HI4PI beam. For these galaxies, we took the sum of the flux within a region that was the optical size of the galaxy at its position and did the same for eight regions surrounding the central position of the galaxy. We divided the summed flux values by the area of the optical size and calculated the standard deviation of the nine values. We then checked that the value was consistent with the values taken from surrounding channels (Res. in Table 1 with data = HI4PI). The exceptions to using this method for large-extent galaxies are the disrupted Sagittarius dSph and Boötes III galaxies. For these galaxies, we quote the limit in the core with the unresolved technique rather than attempt to cover their full extent. The 1σ values from these methods are listed in Table 1, while 5σ is used in the H i mass calculations. The conversion factors from the kelvin σ values to janskys are 0.6 Jy K⁻¹ for HI4PI, 0.11 Jy K⁻¹ for GALFA-H i, and 0.44 Jy K⁻¹ for the GALFA-H i data smoothed to 8' resolution.

Several dwarf galaxies are at velocities that are contaminated by Galactic emission. In most cases, five times the standard deviation calculated with the relevant method above is sufficiently high to represent a limit on the H i gas in the dwarf over the background. In the case of Willman 1, its low heliocentric velocity led to strong Galactic emission and the standard deviation not being suitable for determining the H i mass limit. For this galaxy, we calculated the median flux in 30' regions surrounding the dwarf position in the 10 channels centered on the dwarf velocity and found that twice the median of these values was a conservative value to choose for a limit on what would be a detectable H i mass.

Ultimately, 5σ H i mass limits are quoted in Table 1, and they are calculated using a velocity width of 10 km s⁻¹ and the distance to the dwarf galaxy: M_{H i} (M_⊙) = 2.36 × 10⁵ 5σ (Jy) 10 km s⁻¹ D² (Mpc). We debated whether to use the stellar dispersion values for each galaxy instead of a fixed 10 km s⁻¹, but given the variations on availability and errors, we chose the fixed value. Errors on the mass limits are given based on the distance errors. If the systemic velocity of the dwarf galaxy was not available, we measured the standard deviation every 10 km s⁻¹ for ∣v_LSR∣ < 500 km s⁻¹ in a 30' region at the position of the dwarf (all are unresolved and labeled Unres. no V) and used the average of those values, excluding those channels that were contaminated by Galactic or Magellanic-related emission. The dwarf galaxies without velocities at the time of paper submission were Cetus II, Horologium II, Reticulum III, Pictoris 1, Columba I, Sagittarius II, Indus II, Phoenix 2, Eridanus 3, Indus I, Cetus III, DES J0225+0304, Pictor II, and Virgo I.

Though we checked some of the H i masses for previously detected dwarf galaxies for consistency, we did not remeasure the H i masses and largely adopted values from the references in M12 and adjusted the value if an updated distance was available. The only exceptions to adopting the M12 detections are Sculptor and Fornax. We do not include them because the H i clouds detected are not centered on the optical component of the galaxies, and there is abundant nearby H i that makes the association highly uncertain. Instead, we adopt H i limits at the stellar position of these galaxies. Sculptor formed the vast majority of its stars in the distant past and is in a region where there are multiple high-velocity clouds; therefore, we think it is unlikely that the nearby gas is associated. Fornax has formed stars in the past few gigayears (Weisz et al. 2011), but its velocity makes confusion with Galactic H i emission a huge issue, and the potentially associated extended H i is not centered on the galaxy. The Phoenix dwarf galaxy was considered highly uncertain as well, but with the improved stellar velocities of Kacharov et al. (2017) and the H i clump at the stellar heliocentric velocity and overlapping with the stellar component (St-Germain et al. 1999; Young et al. 2007), we adopt this H i clump as a detection.

In this section, we define the mass models for the Milky Way and M31 that we use to define a Local Group virial radius or "surface." We then describe how we numerically computed orbits for the satellites of the Milky Way that have proper-motion measurements.

3.1. Defining a Local Group "Bounding Surface"

We first define a Local Group coordinate system with an origin at the approximate center of mass of the Milky Way and M31 (computed using models specified below). We define the x-axis of this coordinate system by the line connecting the Galactic center to the center of M31 and place the origin at the center of mass along this axis. Because our models are symmetric about the x-axis, the orientations of the y- and z-axes are arbitrary; with our definition, the z-axis is aligned with (α, δ)_ICRS = (193.327, 48.585)°, and the y-axis is aligned with (α, δ)_ICRS = (102.219, 0.977)°. We assume that the heliocentric distance to M31 is 779 kpc (Conn et al. 2012), the heliocentric distance to the Galactic center is 8.122 kpc (Gravity Collaboration et al. 2018), and the coordinates of the Galactic center are taken from Reid & Brunthaler (2004).

As a simplified representation of the total density distribution of the Local Group, we define a prolate spheroidal surface bounding the Local Group in order to distinguish satellite systems that are within or outside of the Local Group. We first construct a mass model for the total mass distribution of the Milky Way and M31 by representing each galaxy with a spherical Navarro–Frenk–White (NFW; Navarro et al. 1996) density profile such that the total Local Group density is given by the sum of two independent NFW profiles. For the Milky Way, we use virial mass and halo concentration values of M_200,MW = 1.2 × 10¹² M_⊙ and c = 10, motivated by a number of recent mass determinations that use satellite galaxy and globular cluster kinematics to infer the mass of the Galaxy (e.g., Eadie & Jurić 2019; Posti & Helmi 2019; Wegg et al. 2019). For M31, we adopt M_200,MW = 2 × 10¹² M_⊙ (Fardal et al. 2013) and use the same concentration parameter as the Milky Way. The corresponding NFW scale radii are therefore r_s,MW = 22 and r_s,M31 ≈ 27 kpc, respectively. We note that these values are also motivated by and consistent with the measurement of the total Local Group mass M_LG ≈ 3.2 × 10¹² M_⊙ (van der Marel et al. 2012).

To determine the axis ratio of the total Local Group density distribution, q_LG, we compute the xx and yy components of the inertia tensor given by the above mass model and use them to compute the axis ratio q_LG = I_yy /I_xx . We define the Local Group-bounding prolate surface (in Local Group Cartesian coordinates x, y, z) to be

$$\begin{eqnarray}&&{x}^{2}+\displaystyle \frac{{y}^{2}+{z}^{2}}{{q}_{\mathrm{LG}}^{2}}={r}_{e}^{2},\end{eqnarray} \tag{ 1 }$$

where r_e ≈ 733 kpc is defined such that the surface fully encloses the virial sphere around M31. For each satellite galaxy in our sample, we compute the minimum distance between the galaxy and the Local Group-bounding surface. Figure 1 (solid gray ellipse) shows a projection of this bounding surface in both the x_LG, y_LG and x_LG, z_LG planes, along with the positions of all dwarf galaxies in our sample (blue squares for H i–detected galaxies and red upside-down triangles for nondetections) and projections of the virial spheres around the Milky Way and M31 (left and right black circles, respectively).

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We note that there is substantial uncertainty and potentially large (up to a factor of ∼2) biases on the measurements of both the Milky Way and M31 virial masses (e.g., from the presence of the Magellanic Clouds in the case of the Milky Way; Erkal et al. 2020), so this bounding surface should be interpreted qualitatively as a "soft" boundary. To emphasize this point and assess how our conclusions and results are affected by uncertainties on the global properties of the Milky Way and M31 dark matter halos, we also construct two additional Local Group models and bounding surfaces with varied masses for the two galaxies. Since the total Local Group mass is fairly well constrained (van der Marel et al. 2012), we only vary the masses of the Milky Way and M31 to keep the total Local Group mass constant at M_LG = 3.2 × 10¹² M_⊙. As the virial masses of the Milky Way or M31 are not individually well constrained, and the scatter in the existing mass measurements for either galaxy are likely underestimated due to biases in the individual measurements (e.g., Kafle et al. 2018), we choose to vary the Milky Way mass by 25% up and down (adjusting the mass of M31 to keep the Local Group mass fixed). Our models therefore have M_MW = (0.9, 1.2, 1.5) × 10¹² M_⊙, M_M31 = (2.3, 2.0, 1.7) × 10¹² M_⊙, R_vir,MW ≈ (204, 224, 242) kpc, and R_vir,M31 = (279, 266, 252) kpc. The Local Group-bounding surfaces for the two other models are shown in Figure 1 as gray dashed ellipses, and the thickness of the virial circles reflects the range of virial radii for the two galaxies.

In Table 2, we include estimates of the closest distance between each dwarf galaxy and the Local Group surface, D_LG, as estimated in our fiducial mass model. We use the two other mass models to estimate the bounds on the D_LG values for each source and note these as error bars. We note that in some cases (e.g., Pisces II, Pegasus III), the Local Group surface distance in the fiducial mass model is not the central value; this is expected and happens when a dwarf galaxy is near enough to the Local Group-bounding surface such that the change of ellipticity of the surface in the different Local Group mass models causes unintuitive changes.

3.2. A Milky Way Model for Orbit Integration

We numerically integrate orbits for the Milky Way dwarf galaxies that have distances, radial velocities, and proper-motion measurements. For computing these orbits, we use a three-component mass model to represent the Milky Way composed of a Hernquist bulge (Hernquist 1990), Miyamoto–Nagai disk (Miyamoto & Nagai 1975), and spherical NFW halo (Navarro et al. 1997). We fix the parameters of the disk and bulge models to be consistent with the model defined in Bovy (2015). For our fiducial model, we adopt a more massive halo such that the enclosed mass within 250 kpc is M_enc ≈ 1.2 × 10¹² M_⊙ in order to fit a compilation of mass-enclosed measurements (see the gala documentation ⁸ for more information) and be consistent with the model assumed to define the Local Group boundary surface above.

The mass model is implemented in gala (Price-Whelan 2017), which we then also use to numerically integrate the orbits of the dwarfs, treating the satellite galaxies as test particles. We cross-match our catalog of Milky Way satellites with recent proper-motion measurements of a subset of the dwarfs (Fritz et al. 2018; Pace & Li 2019) that make use of astrometry from the Gaia Data Release 2 (Gaia Collaboration et al. 2018b; Lindegren et al. 2018). We compute initial conditions (i.e., at present day) by generating random samples from the error distributions over all kinematic measurements for each of the 38 dwarf galaxies in this subset. For each error sample, we additionally vary the Milky Way halo enclosed mass by sampling a uniform random value between 0.9 and 1.5 × 10¹² M_⊙, following the discussion in Section 3.1, in order to include uncertainty in the Milky Way mass in our derived orbital quantities. ⁹ We transform these samples into a galactocentric reference frame assuming a Sun–Galactic center distance of R_⊙ ≈ 8.122 kpc (Gravity Collaboration et al. 2018) and using solar-motion values from Drimmel & Poggio (2018). We then numerically integrate each orbit sample for each dwarf backward from the present day for 2 Gyr with a time step of 1 Myr using Leapfrog integration. We ignore the gravitational influence of the Magellanic Clouds, which likely affects the orbits of a small subset of the dwarf satellites of the Milky Way (Erkal & Belokurov 2020; Patel et al. 2020). We use these orbits (in particular, pericentric distances and velocities at pericenter) later to determine the minimum Milky Way gas density needed to strip the satellites of their own gas reservoirs. Table 3 includes the orbital parameters for the satellites for which this is done.

Table 3. Median Orbital Parameters and Halo Densities

Galaxy	rperi (kpc)	vperi (km s−1)	nhalo,peri (cm−3) a	rapo (kpc)	vapo (km s−1)	nhalo,apo (cm−3) a
Tucana III b	${3}_{-0.3}^{+0.3}$	536 ± 19	${3.5}_{-3.2}^{+11.1}\times {10}^{-7}$	${37}_{-3}^{+4}$	34 ± 5	${8.8}_{-8.0}^{+28.7}\times {10}^{-5}$
Sagittarius dSph	${15}_{-2}^{+2}$	354 ± 9	${3.8}_{-0.5}^{+0.5}\times {10}^{-4}$	${50}_{-11}^{+19}$	106 ± 20	${4.2}_{-1.3}^{+2.8}\times {10}^{-3}$
Triangulum II	${17}_{-2}^{+3}$	426 ± 29	${5.9}_{-5.3}^{+11.8}\times {10}^{-6}$	${146}_{-36}^{+81}$	50 ± 14	${4.3}_{-3.9}^{+9.6}\times {10}^{-4}$
Segue 1	${18}_{-6}^{+5}$	321 ± 25	${4.8}_{-3.0}^{+4.1}\times {10}^{-5}$	${41}_{-8}^{+17}$	128 ± 18	${3.1}_{-2.0}^{+3.2}\times {10}^{-4}$
Willman 1	${19}_{-8}^{+16}$	306 ± 64	${6.3}_{-2.8}^{+4.8}\times {10}^{-5}$	${43}_{-7}^{+9}$	139 ± 35	${3.1}_{-1.4}^{+3.0}\times {10}^{-4}$
Draco II	${19}_{-1}^{+1}$	378 ± 16	${2.7}_{-2.2}^{+4.2}\times {10}^{-5}$	${97}_{-27}^{+53}$	77 ± 23	${6.5}_{-5.5}^{+15.4}\times {10}^{-4}$
Hydrus I	${25}_{-1}^{+1}$	372 ± 10	${1.9}_{-0.7}^{+0.8}\times {10}^{-5}$	${128}_{-36}^{+85}$	73 ± 25	${5.3}_{-2.8}^{+8.8}\times {10}^{-4}$
Reticulum 2	${25}_{-4}^{+3}$	300 ± 13	${4.4}_{-1.7}^{+2.1}\times {10}^{-5}$	${53}_{-10}^{+14}$	139 ± 19	${2.1}_{-0.9}^{+1.3}\times {10}^{-4}$
Carina II	${27}_{-2}^{+1}$	397 ± 9	${2.8}_{-1.5}^{+2.3}\times {10}^{-5}$	${219}_{-71}^{+176}$	54 ± 19	${1.5}_{-0.9}^{+2.6}\times {10}^{-3}$
Crater II	${27}_{-10}^{+14}$	362 ± 60	${2.0}_{-0.7}^{+1.0}\times {10}^{-5}$	${128}_{-7}^{+9}$	77 ± 20	${4.5}_{-1.8}^{+4.2}\times {10}^{-4}$
Carina III	${28}_{-1}^{+1}$	392 ± 28	${8.7}_{-7.2}^{+15.9}\times {10}^{-5}$	${228}_{-106}^{+306}$	65 ± 37	${3.0}_{-2.6}^{+9.6}\times {10}^{-3}$
Tucana II	${33}_{-9}^{+25}$	358 ± 28	${2.2}_{-1.7}^{+2.9}\times {10}^{-4}$	${160}_{-55}^{+142}$	73 ± 22	${5.1}_{-3.9}^{+9.6}\times {10}^{-3}$
Draco	${33}_{-6}^{+8}$	313 ± 34	${3.1}_{-0.9}^{+1.2}\times {10}^{-4}$	${96}_{-9}^{+11}$	109 ± 9	${2.6}_{-0.7}^{+0.9}\times {10}^{-3}$
Segue II	${33}_{-7}^{+5}$	275 ± 18	${6.2}_{-4.9}^{+11.6}\times {10}^{-5}$	${57}_{-10}^{+22}$	153 ± 22	${2.2}_{-1.7}^{+4.2}\times {10}^{-4}$
Ursa Minor	${34}_{-6}^{+7}$	302 ± 34	${3.7}_{-1.1}^{+1.4}\times {10}^{-4}$	${88}_{-5}^{+6}$	117 ± 10	${2.4}_{-0.7}^{+0.8}\times {10}^{-3}$
Hercules	${36}_{-22}^{+36}$	361 ± 82	${6.9}_{-3.1}^{+4.8}\times {10}^{-5}$	${232}_{-40}^{+138}$	57 ± 28	${2.9}_{-1.5}^{+6.1}\times {10}^{-3}$
Boötes	${38}_{-9}^{+10}$	289 ± 30	${8.9}_{-3.1}^{+4.1}\times {10}^{-5}$	${85}_{-10}^{+20}$	122 ± 15	${5.2}_{-1.8}^{+2.7}\times {10}^{-4}$
Ursa Major II	${38}_{-3}^{+3}$	298 ± 19	${1.3}_{-0.6}^{+0.7}\times {10}^{-4}$	${99}_{-31}^{+61}$	116 ± 34	${8.6}_{-4.6}^{+12.2}\times {10}^{-4}$
Boötes II	${39}_{-2}^{+1}$	396 ± 70	${2.8}_{-2.3}^{+5.3}\times {10}^{-4}$	${537}_{-406}^{+600}$	118 ± 92	${3.0}_{-2.6}^{+14.4}\times {10}^{-3}$
Coma Berenices	${42}_{-2}^{+2}$	284 ± 23	${9.6}_{-3.2}^{+3.9}\times {10}^{-5}$	${96}_{-28}^{+59}$	125 ± 38	${5.0}_{-2.4}^{+6.5}\times {10}^{-4}$
Sculptor	${60}_{-7}^{+8}$	255 ± 19	${4.8}_{-1.3}^{+1.4}\times {10}^{-4}$	${113}_{-12}^{+23}$	134 ± 15	${1.8}_{-0.5}^{+0.7}\times {10}^{-3}$
Sextans	${76}_{-8}^{+8}$	265 ± 14	${3.2}_{-1.0}^{+1.2}\times {10}^{-4}$	${191}_{-53}^{+126}$	106 ± 25	${2.0}_{-0.9}^{+2.0}\times {10}^{-3}$
Horologium 1	${79}_{-22}^{+12}$	245 ± 37	${1.4}_{-1.1}^{+2.1}\times {10}^{-4}$	${129}_{-46}^{+246}$	138 ± 36	${5.1}_{-4.1}^{+10.6}\times {10}^{-4}$
Fornax	${79}_{-27}^{+39}$	242 ± 48	${8.6}_{-3.1}^{+4.2}\times {10}^{-4}$	${152}_{-6}^{+25}$	120 ± 19	${3.5}_{-1.0}^{+1.5}\times {10}^{-3}$
Carina	${81}_{-23}^{+24}$	214 ± 37	${3.2}_{-1.2}^{+2.2}\times {10}^{-4}$	${106}_{-4}^{+15}$	157 ± 19	${6.5}_{-2.4}^{+3.6}\times {10}^{-4}$
Leo I	${101}_{-61}^{+85}$	320 ± 67	${3.0}_{-1.1}^{+0.7}\times {10}^{-4}$	${525}_{-27}^{+204}$	138 ± 55	${1.6}_{-0.8}^{+1.6}\times {10}^{-3}$
Ursa Major	${102}_{-7}^{+6}$	272 ± 49	${2.4}_{-0.9}^{+1.4}\times {10}^{-4}$	${408}_{-256}^{+382}$	121 ± 47	${1.3}_{-0.7}^{+1.6}\times {10}^{-3}$
Aquarius II	${103}_{-7}^{+4}$	433 ± 169	${6.2}_{-5.1}^{+15.1}\times {10}^{-5}$	${1319}_{-1151}^{+915}$	299 ± 203	${1.4}_{-1.2}^{+9.1}\times {10}^{-4}$
Grus I	${116}_{-14}^{+12}$	340 ± 81	${3.1}_{-2.5}^{+6.7}\times {10}^{-5}$	${891}_{-526}^{+506}$	161 ± 116	${1.4}_{-1.2}^{+13.6}\times {10}^{-4}$
Hydra II	${137}_{-32}^{+12}$	456 ± 241	${4.6}_{-4.2}^{+15.0}\times {10}^{-6}$	${1488}_{-1053}^{+1402}$	353 ± 294	${9.3}_{-8.6}^{+112.6}\times {10}^{-6}$
Canes Venatici I	${138}_{-77}^{+46}$	249 ± 64	${3.4}_{-1.3}^{+0.9}\times {10}^{-4}$	${323}_{-64}^{+220}$	106 ± 38	${1.9}_{-0.7}^{+2.4}\times {10}^{-3}$
Leo IV	${154}_{-7}^{+5}$	504 ± 259	${1.5}_{-1.2}^{+4.5}\times {10}^{-5}$	${1799}_{-1269}^{+1349}$	424 ± 291	${2.3}_{-1.9}^{+15.3}\times {10}^{-5}$
Canes Venatici II	${159}_{-49}^{+5}$	316 ± 108	${7.6}_{-4.2}^{+8.8}\times {10}^{-5}$	${786}_{-545}^{+718}$	142 ± 132	${3.5}_{-2.8}^{+10.7}\times {10}^{-4}$
Leo II	${160}_{-115}^{+67}$	226 ± 92	${4.0}_{-2.3}^{+2.4}\times {10}^{-4}$	${239}_{-18}^{+258}$	117 ± 42	${1.5}_{-0.5}^{+2.8}\times {10}^{-3}$
Leo V	${171}_{-6}^{+6}$	543 ± 257	${1.4}_{-1.2}^{+5.2}\times {10}^{-5}$	${1985}_{-1213}^{+1258}$	473 ± 289	${2.0}_{-1.8}^{+13.1}\times {10}^{-5}$
Pisces II	${182}_{-15}^{+15}$	687 ± 375	${2.3}_{-1.9}^{+8.4}\times {10}^{-5}$	${2662}_{-1419}^{+1751}$	636 ± 404	${2.9}_{-2.4}^{+16.8}\times {10}^{-5}$
Eridanus II	${365}_{-17}^{+16}$	955 ± 471	${1.9}_{-1.1}^{+5.0}\times {10}^{-5}$	${3865}_{-1814}^{+2033}$	931 ± 482	${2.0}_{-1.2}^{+6.0}\times {10}^{-5}$
Phoenix c	${417}_{-18}^{+20}$	304 ± 152	${3.3}_{-1.9}^{+6.5}\times {10}^{-4}$	${1306}_{-440}^{+697}$	255 ± 165	${4.5}_{-2.9}^{+17.8}\times {10}^{-4}$

Notes.

^a This is the median halo density required to strip the dwarf at this position. The errors should be multiplied by the same factor of 10 noted for the median density. ^b This galaxy is not included in Figures 6 and 7 because of its very small pericenter. ^c Phoenix is the only galaxy in this list with gas.

A machine-readable version of the table is available.

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The results of the new H i measurements are presented in Table 1 and shown in Figures 4 and 5. All of the newly discovered dwarf galaxies without previous H i measurements are nondetections, and all previous nondetections remain as nondetections. The limits presented here are generally deeper than those previously published. Notable other limits are the deeper limits published for some of the Milky Way dwarf galaxies in Spekkens et al. (2014) and preoptical velocity limits for some of the new ultrafaint dwarf galaxies by Westmeier et al. (2015). To ensure consistency across the complete set of Local Group dwarfs, our limits here depend only on the GALFA-H i and HI4PI data described in Section 2.

Figure 2 shows the relationship between the H i mass limit or H i mass of a dwarf galaxy versus distance from the Milky Way or M31, whichever is closer. The virial radii of the Milky Way and M31, as defined in Section 3, are 224 and 266 kpc, respectively. These virial radii are not presumed to be exact but rather give a guideline to consider the detections and nondetections. A range of possible virial radii for each galaxy is shown by the shaded regions, as described in Section 3.1. The Milky Way and M31 are described individually below, but it can be seen from Figure 2 that the vast majority of the galaxies are nondetections within these radii, and beyond, the vast majority are H i detections. Though we code each dwarf galaxy by whether the Milky Way or M31 is closer, we note that beyond a certain radius, a clear association with either galaxy is uncertain without orbit information. Thus, the majority of the H i detections cannot be directly associated with either the Milky Way or M31. Beyond the virial radii of these galaxies, the detections Leo T, Phoenix, and NGC 6822 could be associated with the Milky Way, and LGS 3, Pegasus dIrr, and IC 1613 could be associated with M31. The other 15 detections beyond the virial radii in this plot are >750 kpc from either galaxy. The average H i mass of the 26 detected dwarf galaxies within 2 Mpc of the Milky Way is 7.5 × 10⁷ M_⊙, and the median H i mass is 7.4 × 10⁶ M_⊙.

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Using the virial radius of 224 kpc for the Milky Way, 53 dwarf galaxies have 5σ limits on their H i mass of <10⁴ M_⊙ within this radius. These are the strongest H i limits due to their proximity and the D² dependency for the H i mass. This dependency also results in the arc of increasing mass limit with increasing distance, since we use survey data that do not vary substantially in sensitivity at a given position. For the Milky Way dwarf nondetections beyond 224 kpc, Leo II has an H i mass limit <10⁴ M_⊙ and, at 233 kpc, is within the range of the possible Milky Way virial radii we consider here, Leo I and Cetus III have H i limits of ∼10⁴ M_⊙ and are at 254 and 251 kpc, and Eridanus II is at 366 kpc with an H i mass limit of <2 × 10⁴ M_⊙. Tucana is a nondetection with a limit of <10⁵ M_⊙ at 887 kpc from the Milky Way and should be considered a field object. The lowest H i mass of a detected dwarf galaxy is above 10⁵ M_⊙, so the results strongly indicate that none of the Milky Way dwarf galaxies with limits are likely to have H i. The only two H i detections within the virial radius of the Milky Way are the Magellanic Clouds, which were not included in GP09 due to their large mass and are currently being stripped of their gas (Putman et al. 2003b; Salem et al. 2015). A radius of 260 kpc would optimize the numbers, with the Milky Way having 56 dwarf galaxies with no gas and two massive dwarf galaxies with gas within this radius. We note that for Sculptor and Fornax, if the offset H i gas was associated (see Section 2), they would be clear outliers in being small dwarf galaxies within 150 kpc of the Milky Way that have H i gas.

The virial radius for M31 from the model in Section 3.1 is 266 kpc. Using this radius as a guide, 27 dwarf galaxies are undetected with limits ≲10⁵ M_⊙. These limits do not vary as much as the Milky Way limits due to the similar distance of these satellites. This count includes M32, And XXV, and And XIX, which have limits slightly above 10⁵ M_⊙ due to some confusion with emission from M31 or the Milky Way. The nondetected dwarf galaxies beyond 266 kpc that can be considered satellites of M31 are And VI at 268 kpc, And XXII at 276 kpc, And XVI at 319 kpc, Perseus I at 348 kpc, And XXVIII at 365 kpc, and And XVIII, which is at a distant 457 kpc. Cetus is a nondetection at 678 kpc from M31 and is in the same category as Tucana, where its large distance makes a direct association with either galaxy uncertain. At close to 2 Mpc from either galaxy, KKR 25 is a nondetection that clearly cannot be associated with either galaxy. The H i detections within the 266 kpc radius for M31 are its dwarf elliptical satellites (NGC 205 and NGC 185) and IC 10. The dwarf ellipticals have smaller amounts of H i relative to their luminosity and dynamical mass compared to other dwarf galaxies (Figures 4 and 5; see also Geha et al. 2006), and IC 10 is at 252 kpc from M31 and has irregular/disturbed H i, similar to LGS 3 at 267 kpc (Hunter et al. 2012; Ashley et al. 2014). This may indicate that the H i gas of these galaxies is beginning to be removed. As can be seen from Figure 2, the outer range of possible M31 virial radii (279 kpc) considered for our Local Group model results in 29 dwarf galaxies within this radius without gas and four with gas.

Figure 3 shows the H i measurements in the context of the dwarf galaxy's distance from a Local Group surface or an approximate virial radius of the Local Group. As outlined in Section 3.1, this surface was defined using the Milky Way and M31 mass and concentration parameters, and variation in this surface was calculated by varying the mass of the Milky Way and M31. The range of distances relative to the variation in the Local Group surface is shown by the line through each symbol, with our fiducial model represented by the position of the symbol and noted in Table 2. It is immediately apparent that, as with the Milky Way and M31, the nondetections are largely located within the Local Group surface (zero on the plot), while those with detected H i are beyond this Local Group surface. There are six detected exceptions within the surface, three of which are the most optically luminous dwarf galaxies in the sample (LMC, SMC, and NGC 205), and two others are in the top 11 most luminous (IC 10 and NGC 185). These detection exceptions are the same as for the Milky Way and M31, with the addition of LGS 3, which is right at the edge of M31's virial radius and very close to the edge of the Local Group surface. The nondetections beyond the Local Group surface are And XVIII, Perseus I, Eridanus II, Leo I, and Leo II. Cetus, Tucana, and the very distant KKR 25 remain odd nondetections at large radii similar to the Milky Way and M31. As can be seen, some dwarf galaxies move inside or outside of the Local Group surface with higher or lower Milky Way and M31 masses. With a 25% lower Milky Way mass but higher M31 mass to obtain the same total Local Group mass, seven (instead of eight) nondetections would be beyond the surface; and with a 25% higher Milky Way mass and correspondingly lower M31 mass, nine nondetections would be beyond the Local Group surface. If the Local Group is more massive than the 3.2 × 10¹² M_⊙ used here, the surface would easily encompass the four nondetections within 50 kpc of the surface. Using our fiducial model as the primary comparison point, we find that the Local Group surface has eight nondetections beyond the surface, and the Milky Way and M31 have a total of 13 nondetections beyond their virial radii. This value of 13 does not change with the changes in the masses of the Milky Way and M31. The Local Group surface therefore encompasses more of the galaxies without gas.

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The completeness of the optical surveys in finding dwarf galaxies at various radii should be considered when assessing the results for both the virial radii of the Milky Way and M31 and the Local Group surface. Most of the dwarf galaxies were discovered in the Sloan Digital Sky Survey, DES, PanSTARRS, or PANDAS (e.g., Willman et al. 2005; Martin et al. 2006; Belokurov et al. 2007; Bechtol et al. 2015; Koposov et al. 2015; Laevens et al. 2015), and though the completeness has been partially assessed for some of these surveys (Walsh et al. 2009; Newton et al. 2018), none cover the complete sky. If one adopts the all-sky survey of Whiting et al. (2007), which is noted to be 77% complete for objects brighter than 25 mag arcsec⁻² in R and one to several arcminutes in size, all of the galaxies within 275 kpc of the Milky Way or Andromeda are devoid of gas. Likewise, galaxies within the Local Group surface or close to the surface used here (e.g., Leo I and Leo II) are devoid of gas. Beyond these distances, things are less clear, with four galaxies detected in H i and three galaxies with H i mass limits. Though the Whiting et al. (2007) sample is very limited in the number of dwarf galaxies, it may suggest that additional faint galaxies without gas will be detected at larger radii.

Figures 4 and 5 examine whether the H i limits for the dwarf galaxies are significant when scaled by their luminosity or dynamical mass. In other words, one might expect the amount of gas in a dwarf galaxy to correlate with the stellar or total mass, and therefore the limits should be lower than the detected galaxies when normalized. The V-band luminosity can be considered as a proxy for the stellar mass, consistent with the stellar mass-to-light ratio of 1 used in M12. While the stellar mass-to-light ratio will vary for different dwarf galaxies, we do not address that here, as these estimates can have many complications and do not span more than an order of magnitude. Figure 4 shows this scaling of the H i content of the dwarf galaxies by V-band luminosity (see Table 2). The scaled H i limits remain significant; however, there are a number of undetected dwarf galaxies around M_{H i} /L_V = 1. The median value of M_{H i} /L_V for the H i–detected dwarf galaxies is 0.93 (blue line in Figure 4), and the average value is 1.2. The undetected dwarf galaxies near these values are dominated by ultrafaint dwarf galaxies with very small stellar populations and some Andromeda dwarf galaxies that are faint but also influenced by slightly higher H i mass limits. All of the undetected Milky Way dwarf galaxies with M_{H i} /L_V ∼ 1 have L_V < 10⁴ L_⊙. Leo T is the best H i–detected comparison point for these galaxies and has an M_{H i} /L_V higher than most of these galaxies. Stronger limits on their H i masses would be useful, but the dwarf galaxies near M_Hi/L_V = 1 have ancient stellar populations and are highly unlikely to have gas that can support star formation.

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The dynamical mass estimates for the dwarf galaxies were calculated within the half-light radius (r_h ) using the same method adopted by M12 in most cases. This involves using the measured stellar velocity dispersions and r_h values and the equation from Walker et al. (2009), ${M}_{\mathrm{dyn}}{(\leqslant {r}_{h})({M}_{\odot })=580{r}_{h}(\mathrm{pc}){\sigma }_{\star }(\mathrm{km}{{\rm{s}}}^{-1})}^{2}$. When values of r_h and σ_⋆ were not in M12, they were added from S19 or the literature, and upper limits were used when necessary (see Table 2). Many of the dwarf galaxies with gas are more distant and do not have a measured stellar velocity dispersion. For these galaxies, we adopt the dispersion calculated from the FWHM of the H i line and calculate the total mass within r_h using ${M}_{\mathrm{dyn}}(\leqslant {r}_{h})=3{r}_{h}{\sigma }_{\mathrm{gas}}^{2}/G$, which can be written ${M}_{\mathrm{dyn}}{(\leqslant {r}_{h})({M}_{\odot })=698{r}_{h}(\mathrm{pc}){\sigma }_{\mathrm{gas}}(\mathrm{km}{{\rm{s}}}^{-1})}^{2}$. The use of the gas dispersion will result in a larger dynamical mass than if the internal stellar velocity dispersion was used. These masses are not designed to be representative of the true total mass of the dwarf galaxies but are used throughout for consistency. Figure 5 shows the results of scaling the H i mass limits and masses by this dynamical mass estimate. It shows that the H i limits are significant when scaled by the total mass interior to r_h , since the vast majority of the detected dwarf galaxies lie above the nondetections. The nondetections with higher M_{H i} /M_dyn values are largely faint Andromeda dwarf galaxies whose H i limits are not as deep. We exclude the Sagittarius dwarf spheroidal galaxy and Boötes III from this plot, since their stellar components are clearly not in equilibrium. For the detections, besides the dwarf ellipticals, Phoenix is somewhat of an outlier in that it is lower than other detections in both this and the M_{H i} /L_V plot. Phoenix is at a velocity of −13 km s⁻¹ and confused with Galactic emission. This location on the plot may indicate more of the gas mixed in with Galactic gas should be associated with the dwarf galaxy, that it is in the process of being stripped, or possibly that the gas is not actually associated with the dwarf.

In this section, we discuss the results in the context of the likely methods of gas removal from Local Group dwarf galaxies. We separate the discussion into the effect of a Milky Way (or M31) gaseous halo medium (Section 5.1) and the potential role of a diffuse Local Group gaseous medium in removing gas from dwarf galaxies (Section 5.2). We briefly discuss other methods to potentially remove gas from dwarf galaxies in Section 5.3.

5.1. Stripping by a Spiral Galaxy Halo Medium

The halo media of the Milky Way and M31 have been thought to play a dominant role in setting the relationship between the gas content of dwarf galaxies and their distance from these large spirals (GP09; Blitz & Robishaw 2000; Spekkens et al. 2014). Numerous dwarf galaxies without gas have now been added within the virial radii of the Milky Way and M31, and this provides additional support for a diffuse halo medium playing an important role in quenching the dwarf galaxies. The few exceptions of H i detections within the virial radii of these galaxies were already known and consistent with being more massive galaxies that can more easily retain their gaseous reservoirs (e.g., Simpson et al. 2018; Garrison-Kimmel et al. 2019). These galaxies also often have gas that is disturbed in nature and may be in the process of being removed. The below discussion focuses on the halo medium of the Milky Way; however, we note that M31 has a halo medium that is likely to be similarly effective at stripping dwarf galaxies (Lehner et al. 2020).

We can freshly examine the required gas density to ram pressure strip Milky Way dwarf galaxies for those that have proper-motion measurements from Gaia (see Table 3; Fritz et al. 2018; Gaia Collaboration et al. 2018a; Simon 2018; Pace & Li 2019). The criterion to ram pressure strip a galaxy of its gas was originally derived by Gunn & Gott (1972). The equation that describes the complete stripping of a homogeneous medium can be written as

$$\begin{eqnarray}&&{n}_{\mathrm{halo}}\sim \displaystyle \frac{{\sigma }^{2}\,{n}_{\mathrm{gas}}}{{v}_{\mathrm{sat}}^{2}}\ ,\end{eqnarray} \tag{ 2 }$$

where n_halo is the ambient gas number density in the halo in cm⁻³, σ is the stellar velocity dispersion of the dwarf, v_sat is the relative motion of the dwarf through the medium, and n_gas is the average gas density of the dwarf in the inner regions. Some comparisons with numerical simulations have found that this equation tends to underestimate the halo density required for stripping. Gatto et al. (2013) completed 2D simulations of the stripping of classical dwarf galaxies and found that the equation gives values that are approximately a factor of 5 too low. Salem et al. (2015) also found a somewhat lower value for the analytical calculation for the stripping effect on the LMC, but the value is within the errors found with the numerical simulation. We leave a more detailed examination of the required halo density to future simulation work but account for this potential large systematic by including large error bars on our calculations.

The internal velocity dispersions (σ) for the Milky Way dwarf galaxies with proper motions are measured in most cases, albeit with large error bars. For Hydra II and Triangulum II, the measured upper limits are used (see Table 2). A dwarf galaxy's dispersion may be similar to its value at earlier times if the dwarf has not been significantly affected by tidal forces and has not grown substantially.

The correct gas density in the inner regions of the dwarf galaxies at the time of stripping is a difficult value to pinpoint. The central densities from model fits to the H i gas profiles depend on the dark matter profiles used and tend to be very high (Faerman et al. 2013; Emerick et al. 2016). On the other hand, star formation can be triggered when gas is compressed as it interacts with a diffuse medium, which can result in a decrease in the gas density. This decrease in the central regions is seen in several galaxies in our sample (e.g., IC 1613 and Sextans A; Hunter et al. 2012). An estimate for the central gas density can be obtained using the current gas distribution in Leo T. Adams & Oosterloo (2018) found a peak column density in the central regions of 4.6 × 10²⁰ cm⁻² in their 117 × 32 pc beam and an H i extent of 400 pc for the dwarf. This results in an average gas density of 0.37 cm⁻³, which is within the range of values used in GP09 and similar to the typical values found in models within the central 100 pc (Emerick et al. 2016). This value is also consistent with what has been found for some larger dwarf galaxies, with central density measurements more typically approaching 10²¹ cm⁻² and extents of a few kiloparsecs (Hunter et al. 2012).

Since a dwarf galaxy will be moving at the highest velocity and most likely through the densest halo medium at perigalacticon, we can assume that this would be the point at which the galaxy would be subject to maximal ram pressure stripping. We therefore compute the minimum halo gas density, n_halo, required to strip the dwarf at perigalacticon using Equation (2), assuming all of the stripping happens at this orbital phase for each dwarf. We note that this is strictly an upper limit on the minimum value, as it does not account for the steady ram pressure stripping from the extended gaseous halo during its orbit. We determine the velocity at perigalacticon for those dwarf galaxies with proper-motion estimates using the Milky Way model and orbit calculations described in Section 3.2. We compute estimates of n_halo and uncertainties on this quantity by Monte Carlo sampling from the error distributions over the velocity dispersion measurements (as compiled from other work) and pericentric velocity for each dwarf galaxy. The error samples in pericentric velocity are computed by sampling from the observed error distribution over distance, proper motion, and radial velocity and numerically integrating the orbits of each dwarf galaxy for the range of Milky Way masses described in Section 3.2. We then compute the median (over samples) value of n_halo (plotted as points in Figure 6) and report upper and lower error bars by computing the 16th and 84th percentiles of our samples over n_halo. We do not account for uncertainty on the gas density of each dwarf because we use a fixed value of n_gas. The orbital values and resulting halo densities are tabulated in Table 3.

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We note that whether all of the dwarf galaxies with measured proper motions are bound remains dependent on the mass of the Milky Way and their potential association with the Magellanic Clouds. In particular, Hyd I, Hor I, Car II, and Car III are likely Magellanic satellites (Simon 2018; Patel et al. 2020) that may be coming in for the first time. In any case, these satellites do not stand out in pericenter or halo density compared to other satellites.

The results of the calculation of the minimum halo density at perigalacticon required to strip the galaxy are shown in Figure 6. For the majority of the dwarf galaxies, the requisite halo gas densities are between 10⁻⁵ and 5 × 10⁻⁴ cm⁻³ out to 200 kpc. The exceptions are two ultrafaint dwarf galaxies that require < 10⁻⁵ cm⁻³ to be stripped and Fornax, which requires ∼10⁻³ cm⁻³ to be stripped. Phoenix and Eridanus have perigalacticons >300 kpc; the former has gas, though it may show some signs of initial stripping, and the latter requires densities of only 2 × 10⁻⁵ cm⁻³ to be stripped. All galaxies with dispersions < 6 km s⁻¹ require halo densities < 1.5 × 10⁻⁴ cm⁻³ to be stripped at perigalacticon. Though direct measures of the density of the Milky Way's halo medium are not available, this type of halo density is easily attainable within 50 kpc with the observational constraints that exist (Sembach et al. 2003; Stanimirović et al. 2006; Hsu et al. 2011; Putman et al. 2011; Miller & Bregman 2015; Salem et al. 2015; Nidever et al. 2019). At larger radii, we do not have solid observational constraints on the halo density of the Milky Way, but absorption line observations indicate that a halo medium is prevalent out to the virial radius of galaxies (Tumlinson et al. 2013; Liang & Chen 2014; Werk et al. 2014). The halo densities required for stripping that are derived here are consistent overall with the values derived from previous related work for some classical dwarfs given the differences in orbital information and central gas densities (e.g., GP09; Gatto et al. 2013).

The derived halo densities can be compared to the halo density with radius values found in simulations, as shown in Figure 7. There are numerous simulations from which the halo gas density can be compared, though most do not plot the data in terms of gas atoms per cubic centimeter verses radius (Putman et al. 2012; Ford et al. 2016; Salem et al. 2016; van de Voort et al. 2019). When the data are made available in this form, the mean is often adopted, which is biased by clumpy higher-density values and results in higher halo densities at a given radius (Kaufmann et al. 2009; Nuza et al. 2014). Indeed, Simons et al. (2020) found that the effect of ram pressure stripping on a satellite galaxy can be highly stochastic owing to the broad dynamic range in density and velocity of the circumgalactic medium (CGM). Therefore, in Figure 7, we choose to show the volume-weighted median and include the 5th and 95th percentiles ¹⁰ to highlight the broad range of densities the dwarf galaxies may experience. The simulations shown are two ENZO cosmological simulations of Milky Way analogs: Joung et al. (2012), which has a total mass of 1.4 × 10¹² M_⊙ at z = 0 within 250 kpc (red), and the Tempest halo from the FOGGIE simulation (Peeples et al. 2019; Zheng et al. 2020), which has a mass of 4.9 × 10¹¹ M_⊙ at z = 0.1 within R₂₀₀ = 161.5 kpc (blue). The Tempest halo from FOGGIE shows lower median values and a broader density spread than the other halo. This difference is likely due to Tempest's lower halo mass and FOGGIE's new refinement scheme that uniformly resolves the CGM of the simulated galaxy in more detail.

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We focus on the Joung et al. (2012) simulation (red) for most of the direct comparison to the dwarf galaxies, given that the Tempest halo from FOGGIE is not as massive as we expect the Milky Way to be. The FOGGIE simulation has higher resolution, and though the spread of properties increases with resolution, the average halo densities remain similar (Peeples et al. 2019; Corlies et al. 2020). The halo densities from Joung et al.'s (2012) simulation are consistent with the majority of the dwarf galaxies being completely stripped as they move through perigalacticon. This is consistent with the simulations of Fillingham et al. (2019) that find that the majority of the galaxies were quenched rapidly on infall. As mentioned previously, Phoenix still has gas, so the fact that it has not passed through a halo medium dense enough to strip it makes sense. There is a group of dwarf galaxies with perigalacticons clustered roughly around 100 kpc that are not as easily stripped by this simulated z = 0 halo medium. Most of these are more massive, with velocity dispersions of 6.6–11.7 km s⁻¹, and Fornax and Carina have longer quenching timescales and more circular orbits consistent with them being less likely to have instantaneous stripping (Fritz et al. 2018; Fillingham et al. 2019). Some of the dwarf galaxies may have been stripped as they passed through an overdense region of halo gas, or star formation triggered by the compression of the gas may have helped to loosen the gas for stripping (Fillingham et al. 2016; Wright et al. 2019; Simons et al. 2020). The halo densities required for stripping are also likely to be more easily reached at earlier times, as cosmological simulations indicate the gaseous surroundings were denser and colder in the past (Fernández et al. 2012; Rahmati et al. 2016). Early stripping is consistent with the fact that these galaxies ceased their star formation at early times.

Given the low densities required for stripping the lowest-mass dwarf galaxies at perigalacticon, we checked the required densities to strip the dwarf galaxies at apogalacticon (see Table 3). This indicates whether these dwarf galaxies could have been stripped early in their orbit or even on entry into the Local Group (see next section). The orbital values at apogalacticon are more uncertain, but generally, gas densities >10⁻⁴ cm⁻³ are required; these are densities that are unlikely to be reached at large radii. The exceptions are a group of small galaxies at apogalacticons >1 Mpc that have required densities of only ∼10⁻⁵ cm⁻³. The level of stripping that a dwarf galaxy experiences throughout its orbit will depend on star formation loosening the gas and the orbital direction of the dwarf galaxy relative to the movement of the halo medium. There are clear indications that the halo medium of a galaxy rotates (Hodges-Kluck et al. 2016; Oppenheimer 2018; Martin et al. 2019; DeFelippis et al. 2020; Simons et al. 2020), and this would potentially make stripping more (or less) effective via the addition of another velocity component between the gaseous medium and the dwarf galaxy.

5.2. Stripping by a Local Group Medium

5.3. Other Methods of Gas Removal

We have placed limits on the gas content of the dwarf galaxies within 2 Mpc of the Milky Way and examined the relationship of gas content with distance from the Milky Way and M31 and relative to their location in the Local Group. The findings can be summarized as follows.

• 1.  
  The H i gas mass limits for the vast majority of the nondetected dwarf galaxies are less than 10⁵ M_⊙ (5σ). This is less than the H i mass of any of the detected dwarf galaxies. This limit improves to <10⁴ M_⊙ for the dwarf galaxies within the virial radius of the Milky Way. Scaling by stellar (i.e., optical luminosity) or total (within the stellar component) mass, the limits are consistent with the gas content of the dwarf galaxies being far suppressed from where they were when they formed their stars.
• 2.  
  The number of dwarf galaxies within 2 Mpc has doubled in the last 10 yr, but there is a distinct lack of new dwarf galaxies with gas. This suggests that galaxies like Leo T are either rare or only at distances where current surveys struggle to detect them (see Tollerud & Peek 2018; DeFelippis et al. 2019). Future surveys capable of detecting low-mass dwarf galaxies at large distances in gas (e.g., WALLABY; Koribalski et al. 2020) and stars (e.g., LSST; Ivezić et al. 2019) will provide important insight into the dwarf population and quenching mechanisms.
• 3.  
  There is a clear relationship that those dwarf galaxies with gas are largely beyond the virial radii of the Milky Way and Andromeda, while nondetections are largely within these radii. This relationship of gas content versus distance is also found when the distance from a Local Group surface, or virial radius, is used. More of the nondetected dwarf galaxies are within the Local Group surface than the virial radii of the Milky Way and M31 (85 ± 1 versus 80, respectively), which may indicate that a Local Group medium plays a role in the gas stripping.
• 4.  
  The relationship of gas content with distance suggests a distance-dependent quenching mechanism such as ram pressure stripping by a diffuse gaseous medium. Using the proper motions available from Gaia, the orbits of 38 dwarf galaxies are calculated, and the minimum required densities for stripping at perigalacticon are typically between 10⁻⁵ and 5 × 10⁻⁴ cm⁻³. Compared to the halo densities found in simulations, 80% of these dwarf galaxies are consistent with being stripped at perigalacticon. Continuous stripping throughout their orbit is likely to play an important additional role in starving and quenching the dwarf galaxies.

The authors acknowledge the GALFA-H i team that enabled much of the data used; useful discussions with Stephanie Tonneson, Andrew Emerick, Juergen Kerp, Kathryn Johnston, Hsiao-Wen Chen, Greg Bryan, Ekta Patel, Tobias Westmeier, Raymond Simons, and Molly Peeples; and useful comments from an anonymous referee. M.E.P. and A.P.W. acknowledge KITP UC-Santa Barbara, where we partially worked on this, which is supported by the National Science Foundation under grant No. NSF PHY-1748958. M.E.P. acknowledges childcare support from Daniel and Elaine Putman during the unique time this paper was submitted.