Pegasus IV: Discovery and Spectroscopic Confirmation of an Ultra-faint Dwarf Galaxy in the Constellation Pegasus

The DELVE survey is an ongoing multicomponent observational campaign seeking to achieve deep, contiguous coverage of the high-Galactic-latitude southern sky in the g, r, i, z bands by combining 126 nights of new observations with existing public archival DECam data. DELVE is split into three main survey components dedicated to studying the resolved stellar substructures and satellite populations of the Milky Way (DELVE-WIDE), the Magellanic Clouds (DELVE-MC), and four nearby galaxies with stellar mass similar to the Magellanic Clouds (DELVE-DEEP). To date, DELVE has taken ∼20,000 new exposures toward this goal, and is expected to finish collecting observations in the 2022B semester. A more detailed description of the DELVE science goals, observing strategy, and progress can be found in Drlica-Wagner et al. (2021).

For this work, we used a new internal photometric catalog for DELVE-WIDE covering nearly the entire sky accessible to DECam with δ_J2000 < +30° and ∣b∣ > 10°, excluding the Dark Energy Survey (DES) footprint. This new catalog will be described in detail in a forthcoming paper (A. Drlica-Wagner et al. 2023, in preparation); we describe the critical components here. We began by selecting all available DELVE and publicly available exposures with exposure times between 30 and 350 s and effective exposure timescale factors t_eff > 0.3 (see Neilsen et al. 2015). After this selection, we were left with a total of ∼40,000 exposures, the largest contributors to which were the Dark Energy Camera Legacy Survey (DECaLS; Dey et al. 2019), the DECam eROSITA Survey (DeROSITAS), ³² and DELVE itself. DELVE-WIDE primarily collects g, i band observations, and the r, z data come primarily from the former two survey programs.

We processed all exposures consistently using the DES Data Management Pipeline (DESDM; Morganson et al. 2018), which reduces and detrends DECam images using custom seasonally averaged bias and flat images, and performs background subtraction. Automated source detection and point-spread-function (PSF) photometry was performed on individual reduced CCD images using SourceExtractor (Bertin & Arnouts 1996) and PSFEx (Bertin 2011). Stellar positions were then calibrated against Gaia Data Release 2 (Gaia Collaboration et al. 2018) using SCAMP (Bertin 2006), and the photometry was calibrated on a CCD-by-CCD basis using zero-points derived from the ATLAS Refcat2 catalog (Tonry et al. 2018) that were transformed into the DECam photometric system (see Appendix B of Drlica-Wagner et al. 2021). Lastly, the resulting calibrated SourceExtractor catalogs for each individual CCD image were merged into a unified multiband object catalog following the procedure introduced in Drlica-Wagner et al. (2015).

Reddening due to interstellar dust was calculated for each object in the resultant catalog from a bilinear interpolation of the maps of Schlegel et al. (1998) with the rescaling from Schlafly & Finkbeiner (2011). Bandpass-specific extinctions were then derived using the coefficients used for DES DR1 (Abbott et al. 2018). Hereafter, we utilize the subscript "0" to denote extinction-corrected magnitudes.

We performed a matched-filter search for old, metal-poor stellar systems in the DELVE-WIDE catalog described above using the simple algorithm ³³ (Bechtol et al. 2015), which has been successfully leveraged to discover more than 20 Milky Way satellites to date. We began by dividing the DELVE-WIDE catalog described in Section 2.1 into HEALPix (Górski et al. 2005) pixels at NSIDE = 32 (∼ 3.4 deg² pixel⁻¹). For each pixel, we selected stars consistent with an old (τ = 12.5 Gyr), metal-poor (Z = 0.0001) PARSEC isochrone (Bressan et al. 2012), which we scanned in distance modulus from 16.0–23.0 mag in intervals of 0.5 mag. Specifically, at each step in the distance modulus grid, we selected all stars with colors consistent with the isochrone locus in color–magnitude space following ${\rm{\Delta }}{(g-r)}_{0}\lt \sqrt{{0.1}^{2}+{\sigma }_{g}^{2}+{\sigma }_{r}^{2}}$. Stars were defined as sources satisfying the criterion

$$\begin{eqnarray*}&&| {\mathtt{SPREAD}}\_{\mathtt{MODEL}}\_{\mathtt{G}}| \lt 0.003+{\mathtt{SPREADERR}}\_{\mathtt{MODEL}}\_{\mathtt{G}},\end{eqnarray*}$$

where the variable SPREAD_MODEL and its associated error, SPREADERR_MODEL, are calculated from a likelihood ratio between the best-fitting local PSF model and a more extended model derived from the same PSF model that is additionally convolved with a circular exponential disk model (Desai et al. 2012). After these selections, the resulting filtered stellar density field was smoothed by a 2' Gaussian kernel, and local density peaks were identified by iteratively raising a density threshold until fewer than 10 distinct peaks remained. Lastly, we computed the Poisson significance of each peak relative to the local background field. Informed by previous searches using simple, we inspected diagnostic plots for all candidates above a significance threshold of 5.5σ.

During visual inspection of the search results produced by simple, we identified a candidate stellar system near (α_J2000, δ_J2000) = (32854, 2662) at a significance of 6.2σ. ³⁴ Within the candidate pool, this system was exceptional because it appeared to display seven stars at g₀ ∼ 20.5 spanning a range of photometric color—a feature indicative of a blue horizontal branch (HB). Querying this candidate's centroid in the SIMBAD database (Wenger et al. 2000) revealed the existence of two RRL variable stars within a radius of 2', both of which were independently identified by the PS1 RRL catalog (Sesar et al. 2017) and the Gaia DR2 variability catalogs (Holl et al. 2018; Clementini et al. 2019).

These identifications strongly merited further investigation of the candidate system. However, the relatively shallow depth of the discovery data was found to be insufficient to draw firm conclusions about the nature and properties of this system. Therefore, we obtained additional g, r, i imaging of the candidate system during regular DELVE observing and in DECam engineering time in 2021 August. These newer observations consisted of 333 s exposures centered on the candidate, improving the depth by ∼0.4 mag in each band compared to the discovery data. These deeper exposures were then incorporated into a newer iteration of the DELVE catalog (prepared identically to the catalog described in Section 2.1), and this newer catalog was used for all analyses and figures in the following sections.

In Figure 1, we present diagnostic plots for the candidate stellar system similar to those generated for each overdensity identified by simple . These include the smoothed distribution of isochrone-filtered stars and galaxies (leftmost and center-left panels, respectively), a background-subtracted Hess diagram (center-right panel), and a radial profile for the system (rightmost panel), including the best-fit Plummer (1911) model derived in Section 3.

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Our analyses described in the following sections strongly suggest that this system is an ultra-faint dwarf galaxy, rather than a star cluster. Therefore, following the historical naming convention for confirmed dwarf galaxy satellites of the Milky Way, we refer to the system as Pegasus IV throughout this work.

To confirm that Pegasus IV is a bound stellar system, and to determine its kinematic and dynamical properties, we observed the system with the 6.5 m Magellan-Baade Telescope and IMACS (Dressler et al. 2011) on a two-night observing run spanning 2021 September 12–13. Following previous studies of ultra-faint dwarf galaxies using IMACS, we used the instrument's f/4 camera and the 1200 ℓ mm⁻¹ grating blazed at 9000 Å (e.g., Simon et al. 2017). The resulting spectra spanned a wavelength range of ∼7500–9000 Å at R ∼ 11,000, sufficient for precise velocity and metallicity measurements from the CaT absorption feature centered at roughly 8500 Å.

We observed a single multislit mask centered on the system, which featured 32 07 × 5'' slits. Targets were chosen in the following order. First, we selected red giant branch (RGB) and HB stars consistent with a Dotter (2016) isochrone with age τ = 12.5 Gyr and metallicity [Fe/H] = −2.3 in our DECam photometry, informed by past studies of ultra-faint dwarf galaxies. We then added bright stars that we identified as possible members on the basis of a preliminary mixture model analysis of their proper motions in Gaia EDR3 (see Section 4.7). Lastly, to fill remaining available space on the slitmask, we added several stars from Gaia that lacked DECam photometry.

Due to the northern decl. of Pegasus IV (δ₂₀₀₀ ∼ +27°) and the southern latitude of Las Campanas Observatory, we were only able to observe Pegasus IV at airmass ≲1.8 with Magellan/IMACS for a little over an hour on each night. On each night, we collected two science exposures (1800 + 2400 s), followed by (Kr, Ar, Ne, He) arc lamp calibration frames and flat frames. The typical seeing for these observations was 11 on September 12 and 075 on September 13.

We reduced the IMACS spectroscopic observations following the procedure described by Simon et al. (2017). In brief, this process first involved using the Cosmos reduction pipeline (Dressler et al. 2011; Oemler et al. 2017) to map slits on the IMACS detector plane and achieve a preliminary wavelength solution based on the arc lamp data. Then, a modified version of the DEEP2 data reduction pipeline (Cooper et al. 2012; Newman et al. 2013) was used to extract and calibrate the one-dimensional spectrum for each star. We then combined the spectra from the four exposures using inverse-variance weighting.

We measured stellar radial velocities from the IMACS spectra following the method introduced in Li et al. (2017). This method involves fitting the reduced spectrum of each star with velocity templates by shifting the template through a range of velocities to find the velocity v_obs that maximizes the likelihood

$$\begin{eqnarray}&&{ \mathcal L }=-\displaystyle \frac{1}{2}\sum _{\lambda ={\lambda }_{1}}^{{\lambda }_{2}}\displaystyle \frac{{\left[{f}_{\mathrm{spec}}(\lambda )-{f}_{\mathrm{temp}}\left(\lambda \left(1+\tfrac{{v}_{\mathrm{obs}}}{c}\right)\right)\right]}^{2}}{{\sigma }_{\mathrm{spec}}^{2}}.\end{eqnarray} \tag{ 1 }$$

Here, f_spec(λ) and ${\sigma }_{\mathrm{spec}}^{2}(\lambda )$ represent a normalized spectrum and its corresponding variance, and f_temp represents a normalized velocity template spectrum. Because we measured velocities specifically from the CaT absorption feature, we set the wavelength bounds of the spectral fit to be λ₁ = 8450 Å and λ₂ = 8685 Å. All of our IMACS spectra were fit with three velocity templates: HD122563, a very metal-poor RGB star; HD26297, a more metal-rich RGB star; and HD161817, a blue HB star. We report the velocity measurement from the template that produced the largest likelihood at the best-fit velocity.

For each spectrum-template combination, we ran the MCMC sampler implemented by emcee to sample the likelihood function above. To ensure robust sampling, we used 25 walkers each taking 2000 steps, with the first 500 steps for each walker discarded as burn-in. Then, for each star, we took the median and the standard deviation (after 5σ clipping) of the velocity posterior distribution for the best-fit template as the measured velocity v_obs and velocity error ${\sigma }_{{v}_{\mathrm{obs}}}$, respectively.

We next applied a telluric correction to this measured velocity v_obs to account for the miscentering of stars within slits, which can lead to small (<10 km s⁻¹) offsets in the measured velocities of stars (see e.g., Sohn et al. 2007). To derive the correction for each spectrum, we reran the identical template-fitting MCMC procedure described above except with a telluric template, setting λ₁ = 7550 Å and λ₂ = 7700 Å. The median and standard deviation of the resulting posterior distribution then provided the magnitude of the telluric correction v_tell and its associated variance ${\sigma }_{{v}_{\mathrm{tell}}}^{2}$.

The corrected velocity of each star, v, was calculated as v = v_obs−v_tell with an associated uncertainty of ${\sigma }_{{v}_{\mathrm{stat}}}\,=\sqrt{{\sigma }_{{v}_{\mathrm{obs}}}^{2}+{\sigma }_{{v}_{\mathrm{tell}}}^{2}}$. The error ${\sigma }_{{v}_{\mathrm{stat}}}$ is purely statistical in nature, and is directly correlated with the signal-to-noise ratio (S/N) of each individual spectrum. Informed by previous studies that considered the repeatability of IMACS velocities between successive nights (e.g., Simon et al. 2017; Li et al. 2018), we also added a 1.0 km s⁻¹ systematic error term in quadrature to each velocity measurement error.

In summary, the above steps resulted in velocities v for each star, each with a single associated uncertainty. These velocities were then transformed into the heliocentric frame. For the rest of this work, we denote the resulting heliocentric velocities as v_hel. In total, we were able to measure reliable velocities for 23 unique stars at S/N > 3.

We measured the metallicity of RGB member stars in Pegasus IV through the equivalent widths (EWs) of the CaT lines. We modeled each of the three CaT lines for each star with a Gaussian-plus-Lorentzian profile (e.g., Hendricks et al. 2014; Simon et al. 2015), and converted their summed EWs to [Fe/H] metallicities using the calibration relation from Carrera et al. (2013). This relation requires an absolute V-band magnitude for each star, and thus we first converted from the DELVE g, r-band photometry to this system using the relation provided in Bechtol et al. (2015), and then subtracted the distance modulus derived from the ugali fit (Section 3). The resulting error on the metallicity for each star was fully propagated from a combination of four sources: (1) uncertainty in the EW measurements, including a 0.2 Å systematic uncertainty floor (Li et al. 2018), (2) uncertainties in the coefficients from the Carrera et al. (2013) relation, (3) uncertainties in the DELVE photometry, and (4) uncertainty associated with the distance modulus from ugali. Specifically, we Monte Carlo sampled the posterior distribution associated with each of the above sources of uncertainty, applied the appropriate formulae to calculate [Fe/H] for each set of samples, and then used the median and 16th/84th percentiles of the resultant [Fe/H] samples to define our measurement and 1σ confidence interval for each star. The first of these sources of error is dominant for all but the brightest member star (see Section 4.4). ³⁶

In general, accurate CaT EW measurements require higher S/N than accurate velocity measurements. Visual inspection of the spectra for stars in our sample revealed that the CaT fits for stars with low S/N were of poor quality, and thus we opted to impose an S/N > 5 cut for metallicity measurements. In total, we measured metallicities for 11 stars above this threshold.

From the 23 spectra with S/N > 3 for which we measured velocities, we identified a clear clustering of 12 stars with radial velocities −300 ≲ v_hel ≲ −250 km s⁻¹, including nine within the narrower range of −282 km s⁻¹ ≲ v_hel ≲ −262 km s⁻¹ (see top right panel of Figure 2). These 12 stars were separated in velocity from all other measured stars with S/N > 3 by a gap of >100 km s⁻¹, and were all located within $4^{\prime} \ (\sim 2.5{r}_{h})$ of our derived centroid for Pegasus IV. We summarize the key properties of these 12 stars in Table 2.

Table 2. Properties of Pegasus IV Spectroscopic Member Candidates, Ordered by Decreasing IMACS Spectrum S/N

Gaia EDR3 SourceID	R.A.	Decl.	g0	r0	S/N	vhel	[Fe/H]	Comment
	(deg)	(deg)	(mag)	(mag)		(km s−1)	(dex)
1796890071833434112	328.52536	26.62094	18.01	17.07	52.4	−271.9 ± 1.0	$-{2.00}_{-0.11}^{+0.11}$	RGB
1796887219975171328	328.59498	26.63219	19.77	19.16	10.6	−278.1 ± 1.3	$-{2.82}_{-0.41}^{+0.34}$	RGB
1796888907896667520	328.48786	26.63070	⋯ a	⋯ a	7.9	−277.9 ± 1.7	$-{2.85}_{-0.40}^{+0.32}$	RGB
1796888834882857216	328.49587	26.62398	20.05	19.54	7.8	−257.9 ± 2.0	$-{2.77}_{-0.37}^{+0.31}$ b	Binary/nonmember?
1796890381071133568	328.52503	26.64546	20.44	19.87	5.5	−269.2 ± 2.2	$-{3.29}_{-0.51}^{+0.31}$	RGB
1796890071833414784	328.53716	26.61518	20.13	19.87	4.7	−265.1 ± 1.8	⋯	HB; variable?
1796886807658193536	328.59171	26.58421	20.54	20.01	4.6	−295.6 ± 2.5	⋯	Nonmember?
1796887082536156928	328.56077	26.60775	20.51	20.20	3.6	−271.8 ± 3.7	⋯	HB; RRL
1796891171345139456	328.49033	26.65611	20.80	20.29	3.5	−271.2 ± 4.2	⋯	RGB
1796887048176397952	328.56543	26.59852	20.83	20.25	3.5	−288.1 ± 2.8	⋯	Nonmember?
1796887151255658752	328.56076	26.61745	20.98	20.40	3.3	−270.7 ± 3.2	⋯	RGB
1796890071833423872	328.52913	26.61916	21.00	20.49	3.2	−271.75 ± 3.7	⋯	RGB

Notes. R.A. and decl. coordinates are taken from Gaia EDR3. The g, r-band photometric measurements are taken from DELVE (with one exception; see below), and correspond to AB magnitudes in the DECam photometric system. The reported S/Ns refer to the IMACS spectroscopic data. The velocity uncertainties reported here include the 1.0 km s⁻¹ systematic uncertainty discussed in Section 4.2. We provide a similar table reporting the properties of the spectroscopically observed nonmembers in Appendix B.

^a This star was not in the DELVE photometric catalog, as it was obscured by a charge-bleed artifact caused by a nearby bright star. For the purpose of deriving this star's [Fe/H] metallicity, we instead calculated a V-band magnitude for this star using the Gaia EDR3 photometry, adopting the transformation relation provided in Table 5.7 of the official Gaia EDR3 documentation; we assumed a conservative error of ±0.1 mag. This magnitude was then used in the Carrera et al. (2013) relation. ^b This metallicity assumes that the star is a member of Pegasus IV, and therefore that its distance is ∼90 kpc.

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To assess which stars among this sample of 12 were plausible Pegasus IV members as opposed to Milky Way contaminants, we subjectively inspected these stars' proper motions from Gaia EDR3, locations in color–magnitude space from the DELVE photometry, and heliocentric velocities and metallicities from the IMACS spectroscopy (where possible). We found that all 12 stars displayed self-consistent proper motions (within 1σ–2σ) and were photometrically consistent with an old, metal-poor isochrone (see bottom panels of Figure 2). Thus, we found no reason to reject any stars as members on the basis of color or proper motion information.

The velocities of these 12 stars appeared to show a considerable spread, ranging from −258 km s⁻¹ ≲ v_hel ≲ −296 km s⁻¹. As can be seen in Figure 3, nine of these stars lay within 10 km s⁻¹ of the apparent mode near v_hel ∼ −272 km s⁻¹. The remaining three stars fell significantly outside of this range, lying at v_hel ∼ [−258, −288, −296] km s⁻¹. Even if Pegasus IV truly exhibits a large velocity dispersion, these stars' separation from the peak of the observed velocity distribution suggested that they are either nonmembers or are binary star members of Pegasus IV that were observed at an orbital phase that places them far from their center-of-mass velocity. ³⁷

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The first of the aforementioned velocity outliers (at v_hel ∼ −257 km s⁻¹) was found to have [Fe/H] = $-{2.77}_{-0.37}^{+0.31}$ dex—consistent with the mean metallicity of this system (see Section 4.6). In general, isolated halo stars at this metallicity are relatively rare (e.g., Schörck et al. 2009; Youakim et al. 2020). Thus, we presume that this star is a candidate binary member star of Pegasus IV, rather than a nonmember, but emphasize that this assumption has a significant impact on the measured velocity dispersion (see Section 4.5). We further caution that the metallicity estimate for this star assumes that the star is located at the distance of Pegasus IV, and will be underestimated should the star prove to be a foreground contaminant. In contrast to the case of the velocity dispersion, however, this star has only a small effect on our estimation of Pegasus IV's metallicity dispersion (see Section 4.6).

For the remaining two velocity outliers, both of which appeared to lie along the RGB of our best-fit isochrone, we were unable to confidently distinguish whether these stars are binary members or foreground nonmembers in the absence of multiepoch velocity measurements. To assess the likelihood that these two stars could be Milky Way foreground contaminants, we ran a simulation using the web interface to the Besançon Galactic model (Robin et al. 2003). We first queried the model to produce a catalog of stellar magnitudes and kinematic measurements for simulated stars within a 1 deg² region centered on Pegasus IV. We then transformed the resultant magnitudes from the SDSS photometric system to the DECam photometric system using the equations provided by Drlica-Wagner et al. (2018). Then, we computed the expected surface density of Milky Way stars within a radius of $r\lt 30^{\prime}$ that were consistent with the RGB of our target selection isochrone, had heliocentric radial velocities −300 km s⁻¹ ≤ v_hel ≤ −250 km s⁻¹, and had small proper motions (∣μ∣ < 4 mas yr⁻¹ in each direction). After multiplying this surface density by the area of the region that the IMACS slitmask covered (∼100 arcmin²), we found that ∼1 foreground star is expected in our spectroscopic sample within this velocity range. Our observation of two stars with outlying radial velocities is slightly inconsistent with this prediction, potentially suggesting that one or both of these stars is a binary member of Pegasus IV. We reiterate that the membership status of these two stars remains highly uncertain.

Lastly, to assess whether the brightest star was indeed a member star despite its relatively high metallicity ([Fe/H] = −2.00 ± 0.11 dex; first row of Table 2), we measured the equivalent width of its Mg i λ8807 Å absorption line. As described by Battaglia & Starkenburg (2012), this line can be used in conjunction with the CaT to discriminate between foreground Milky Way contaminants (primarily main-sequence stars) and dwarf galaxy members (red giants). Fitting the Mg I line with a Gaussian profile, we calculated the equivalent width to be 0.16 ± 0.02 Å (statistical error only). Given the stars CaT equivalent width of 5.1 ± 0.1 ± 0.2 Å , this confidently places the star in the red giant regime defined by Equation (1) of Battaglia & Starkenburg (2012), and thus we concluded that it is very likely that this star is a true RGB member of Pegasus IV. ³⁸

In summary, we identified nine clear spectroscopic member stars, in addition to one candidate binary member and two potential members with considerably uncertain status. Of the nine clear members, seven are RGB stars, and two appear to lie on the HB. The clear members are shown as blue triangles in Figure 2, while the potential members with uncertain status are shown as black triangles. One of the two spectroscopically observed HB stars is classified as an RRL-type variable in the PS1 and Gaia RRL catalogs (see Section 5.5), and the other appears to show some signs of photometric variability in our data. We include all 12 of these stars in Table 2 as candidate members, but include a comment in the final column to highlight each of the uncertain cases.

To constrain the systemic velocity (v_hel) and velocity dispersion (σ_v ) of Pegasus IV, we sampled the two-parameter Gaussian likelihood function defined by Equation (8) of Walker et al. (2006) using emcee. We applied a uniform prior on v_hel with a range of [−250, −300] km s⁻¹, and a uniform prior on $\mathrm{log}({\sigma }_{v})$ with a range of [−2, 2]. For our primary kinematic measurements, we included only the seven clear (nonoutlier) RGB member stars described in Section 4.4. We excluded the two candidate variable stars on the HB from our kinematic sample, since the pulsation of variable stars causes their apparent velocities to vary over time.

Using these seven stars, and applying the priors described above, we measured Pegasus IV's systemic velocity to be ${v}_{\mathrm{hel}}=-{273.6}_{-1.5}^{+1.6}\ \mathrm{km}\ {{\rm{s}}}^{-1}$ with a velocity dispersion of ${\sigma }_{v}\,={3.3}_{-1.1}^{+1.7}\ \mathrm{km}\ {{\rm{s}}}^{-1}$. The resulting posterior probability distributions from the MCMC sampling are shown in the left panel of Figure 4, and the best-fit model is depicted in black over the velocity histogram in Figure 3. To assess the impact of our prior on this measurement, we also explored adopting a flat prior on σ_v , rather than $\mathrm{log}({\sigma }_{v})$. Holding all else constant, this change of prior resulted in changes to the systemic velocity and velocity dispersion that were significantly smaller than the quoted errors of our primary measurements.

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Our measured velocity dispersion of ${\sigma }_{v}={3.3}_{-1.1}^{+1.7}\ \mathrm{km}\ {{\rm{s}}}^{-1}$ is clearly nonzero, implying that we confidently resolved the internal dynamics of Pegasus IV. However, the value of this dispersion was found to be sensitive to the exact member used in our velocity dispersion fit. In particular, we observed that including the metal-poor outlier at v_hel ∼ −257 km s⁻¹ (while retaining our default priors) raised the velocity dispersion to ${\sigma }_{{v}_{\mathrm{hel}}}={6.0}_{-1.3}^{+2.0}\ \mathrm{km}\ {{\rm{s}}}^{-1}$, consistent within ≲1.5σ of the measured dispersion from our nominal seven-star sample. Similarly, including all 12 candidate member stars would raise the velocity dispersion to ${\sigma }_{v}={10.0}_{-1.9}^{+2.8}\ \mathrm{km}\ {{\rm{s}}}^{-1}$ (after relaxing the $\mathrm{log}({\sigma }_{v})$ prior to [−3, 3]). Given that adopting these alternate member samples only increased the resulting velocity dispersion, our primary results derived from the seven-star sample can be considered the most conservative estimate of the dark matter content of Pegasus IV. This ensures that our ultimate conclusion that Pegasus IV is a dark-matter-dominated dwarf galaxy (see Section 5.1) is insensitive to assumptions about the nature of these apparent velocity outliers.

Therefore, under the assumption that Pegasus IV is a dispersion-supported system in dynamical equilibrium, we proceeded to estimate the system's dynamical mass using the mass estimator introduced in Equation (2) of Wolf et al. (2010):

$$\begin{eqnarray}&&{M}_{1/2}\approx 930\left(\displaystyle \frac{{\sigma }_{v}^{2}}{{\mathrm{km}}^{2}\,{{\rm{s}}}^{-2}}\right)\left(\displaystyle \frac{{r}_{1/2}}{\mathrm{pc}}\right){M}_{\odot }.\end{eqnarray} \tag{ 2 }$$

Using our measured dispersion from the nominal seven-star sample and the half-light radius from Section 3, we found Pegasus IV's enclosed mass within r_1/2 to be ${4.0}_{-2.3}^{+5.1}\times {10}^{5}{M}_{\odot }$. The mass-to-light ratio within one half-light radius is therefore ${M}_{1/2}/{L}_{V,1/2}={166}_{-99}^{+224}{M}_{\odot }/{L}_{\odot }$.

To measure Pegasus IV's mean metallicity ([Fe/H]_spec) and metallicity dispersion (σ_[Fe/H]), we applied a simple Gaussian likelihood model that was nearly identical to the model used for the velocity and velocity dispersion. We adopted a uniform prior on $\mathrm{log}({\sigma }_{[\mathrm{Fe}/{\rm{H}}]})$ with a range of [−2, 2], and again performed MCMC sampling using emcee. By default, we opted only to use the five stars with S/N > 5 (including the binary candidate at v_hel ∼ −257 km s⁻¹). ³⁹ For these stars, we found ${[\mathrm{Fe}/{\rm{H}}]}_{\mathrm{spec}}=-{2.63}_{-0.30}^{+0.26}$ dex and metallicity dispersion ${\sigma }_{[\mathrm{Fe}/{\rm{H}}]}={0.47}_{-0.18}^{+0.30}$ dex. The resulting posterior probability distributions are shown in the right-hand panel of Figure 4.

Given the small sample of stars with S/N > 5 (five in total), we performed a jackknife test (Efron 1982) to assess the robustness of our measured metallicity and metallicity dispersion. We removed individual stars, one at a time, and reran the MCMC sampling. After doing so, we found that the brightest, most metal-rich star (first row of Table 2) had a particularly strong influence on the metallicity and metallicity dispersion. Removing this star resulted in a more metal-poor systemic mean metallicity of [Fe/H] = −2.91 ± 0.20 dex and upper limit on the metallicity dispersion of 0.1 dex (at 84% confidence), in ∼2σ tension with the five-star measurement. By contrast, removing each of the three stars at [Fe/H] ∼ −2.82 dex (second, third, and fourth rows of Table 2) minimally affected the mean metallicity and minorly increased the measured dispersion (at the level of ∼0.07 dex, well within the uncertainties on the five-star sample dispersion). Lastly, removing the most metal-poor star (fifth row of Table 2, [Fe/H] ∼ −3.3 dex) increased the mean metallicity to [Fe/H] = $-{2.45}_{-0.35}^{+0.21}$ dex and resulted in a slightly smaller dispersion of ${\sigma }_{[\mathrm{Fe}/{\rm{H}}]}={0.40}_{-0.35}^{+0.21}$ dex. Both of these are consistent within 1σ of the five-star result.

These results, in aggregate, suggest that Pegasus IV is a metal-poor stellar system with a tentative detection of a nonzero metallicity dispersion. The magnitude of the dispersion is highly contingent on the membership of the brightest star. Since our measurement of the Mg I λ8807 Å line for this star gives no reason to doubt its membership, we opt to report the metallicity and corresponding dispersion from the five-star sample, namely, ${[\mathrm{Fe}/{\rm{H}}]}_{\mathrm{spec}}=-{2.63}_{-0.30}^{+0.26}$ dex and ${\sigma }_{[\mathrm{Fe}/{\rm{H}}]}={0.47}_{-0.18}^{+0.30}$ dex, but we emphasize that the value of our measured dispersion is tentative and should be interpreted cautiously due to the small sample size. We also note that this includes the star with a velocity of v ∼ 257 km s⁻¹, which we assumed to be a true member, but excluded in our kinematic analysis given the likelihood that this star is an unresolved binary.

Regardless of the input stellar sample, the mean metallicity of Pegasus IV places it among the most metal-poor ultra-faint dwarfs known, which include Reticulum III, Bootes II, Tucana II, Horologium I, Draco II, Reticulum II, and Grus I, the metallicities of which range from −2.81 < [Fe/H] < −2.62 dex. Our measurement suggests that Pegasus IV is slightly more metal-poor than other dwarf galaxies of similar absolute magnitude (see the right panel of Figure 5), but this difference does not appear to be statistically significant.

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Our measured metallicity dispersion (${\sigma }_{[\mathrm{Fe}/{\rm{H}}]}={0.47}_{-0.18}^{+0.30}$ dex), while relatively uncertain, is comparable to the dispersions observed in other ultra-faint dwarf galaxies at similar absolute magnitude, i.e., Columba I, Coma Berenices I, Leo V, Pisces II, and Ursa Major II, which have σ_[Fe/H] = 0.71, 0.43, 0.30, 0.48, and 0.67 dex, respectively (Kirby et al. 2015; Fritz et al. 2019; Simon 2019; Jenkins et al. 2021). Pegasus IV's metallicity dispersion can also be compared to the intrinsic iron abundance spreads observed in the Milky Way's (bright) globular cluster population, which Bailin (2019) found to have a median metallicity dispersion of σ_[Fe/H]=0.045 dex across a sample of 55 clusters with high-resolution spectra.

We computed the systemic proper motion of Pegasus IV using the precise astrometry provided by Gaia EDR3 (Gaia Collaboration et al. 2021). We analyzed three different proper motion models to measure the systemic proper motion. The first was a mixture model composed of dwarf and Milky Way components and utilizes spatial position and proper motion (Pace & Li 2019; Pace et al. 2022). This model was run prior to the acquisition of both deeper photometry and spectroscopy, and its results informed our spectroscopic target selection. It was run with preliminary spatial parameters and only used stars with DECam photometry. With this model, we found μ_α* = 0.39 ± 0.17 mas yr⁻¹ and μ_δ = 0.01 ± 0.18 mas yr⁻¹. The model also reports the number of probable members with proper motion measurements, which was found to be N = 12.3 ± 1.4 stars. Due to the color–magnitude selection window, this model missed the brightest member, which explains the worse precision compared to the following models.

The second proper motion model was similar to the first, but used only a fixed sample of spectroscopic members as input. We used a multivariate Gaussian distribution to model the dwarf, and we sampled the posterior probability using emcee. With this model, we found μ_α* = 0.33 ± 0.07 mas yr⁻¹, μ_δ = −0.21 ± 0.08 mas yr⁻¹, assuming a fixed sample of N = 9 stars (consisting of the seven stars used for the dynamical analysis, in addition to the two spectroscopically observed variable star candidates).

The third model used a similar mixture model, but built on those of Pace & Li (2019) and Pace et al. (2022) by incorporating spectroscopic information. We preassigned the membership of stars with spectroscopy, which assists in determining both the dwarf and Milky Way proper motion distributions. We did not exclude stars missing DECam photometry and instead applied a loose Gaia color–magnitude selection for these stars. With this model, we found μ_α* = 0.33 ± 0.07 mas yr⁻¹ and μ_δ = −0.22 ± 0.08 mas yr⁻¹, and N = 13.1 ± 0.6. The proper motion from this model is almost identical to the spectroscopic member-only results from the second model, likely because the additional members are generally faint (mostly HB stars) and do not significantly influence the measurement. We note that the majority of the systemic proper motion precision comes from the brightest member. The proper motion error of this star is ${\sigma }_{{\mu }_{\alpha * }}\,=0.08\ \mathrm{mas}\,{\mathrm{yr}}^{-1}$ (similar to the systemic proper motion error) and its inclusion decreases the systemic proper motion error by ∼40%. We opted to use the systemic proper motion derived from the spectroscopic members as our preferred measurement for further analysis of Pegasus IV's kinematics, since this measurement is least likely to be biased by contaminant stars. We do note, however, that differences have been observed between dwarf galaxy proper motions derived from spectroscopic samples and those derived without (e.g., Massari & Helmi 2018).

Recent discoveries of ultra-faint Milky Way satellites have broadly consisted of two classes of objects: dark-matter-dominated dwarf galaxies and likely baryon-dominated halo star clusters. We find that Pegasus IV is significantly more consistent with the former class of objects on the basis of its size, mass-to-light ratio, and metallicity dispersion. Specifically, Pegasus IV's half-light radius is larger than the population of known globular clusters (see the left panel of Figure 5). More conclusively, Pegasus IV's large mass-to-light ratio (${M}_{1/2}/{L}_{V,1/2}={167}_{-99}^{+224}\,{M}_{\odot }/{L}_{\odot }$) is inconsistent with the known population of halo star clusters, which typically exhibit mass-to-light ratios of ∼1–3 M_⊙/L_⊙ (e.g., Dalgleish et al. 2020). Lastly, the system's tentatively resolved metallicity dispersion suggests that it has undergone multiple generations of star formation and/or that its gravitational potential well is deep enough to have retained supernova ejecta, both of which are indicative of a dark-matter-dominated dwarf galaxy (e.g., Willman & Strader 2012).

We note that the conclusion that Pegasus IV is an ultra-faint dwarf galaxy could be further tested in the future through higher-resolution spectroscopic observations of its bright member stars. Such spectra would allow for measurements of the galaxy's α-element and neutron-capture element abundances, both of which can independently offer further insight into the classification of this system (e.g., Ji et al. 2019). Alternately, deeper medium-resolution spectra could provide iron abundances for a large sample of stars, allowing for a more robust measurement of the system's metallicity dispersion.

To determine Pegasus IV's orbital properties, we integrated 500 realizations of its orbit using the gala Python package (Price-Whelan 2017). For each realization, we determined Pegasus IV's initial conditions {α_J2000, δ_J2000, D_⊙, μ_α*, μ_δ , v_⊙} by sampling from the error distributions of its observed position and kinematics (Table 1), which we approximated as Gaussian for all parameters. We then rewound Pegasus IV's orbit back in time for 5 Gyr in the presence of gala's default Milky Way model, which includes a spherical nucleus and bulge, a Miyamoto–Nagai disk (Miyamoto & Nagai 1975), and a spherical Navarro–Frenk–White (NFW) dark matter halo (Navarro et al. 1996).

At the conclusion of each integration, we recorded gala's estimate for Pegasus IV's apocenter (r_apo), pericenter (r_peri), eccentricity (e), orbital angular momentum perpendicular to the Galactic disk (L_z ), and total energy (E). From the median and 16th/84th percentiles of the distributions for these quantities across the 500 realizations, we find

• 1.  
  ${r}_{\mathrm{apo}}={94}_{-7}^{+8}\,\mathrm{kpc}$ ${r}_{\mathrm{peri}}={32}_{-14}^{+18}\,\mathrm{kpc}$
• 2.  
  $e={0.49}_{-0.16}^{+0.17}$
• 3.  
  ${L}_{z}={6.3}_{-2.6}^{+2.9}\,{\mathrm{kpc}}^{2}\,{\mathrm{Myr}}^{-1}$
• 4.  
  $E=-{0.049}_{-0.004}^{+0.005}\,{\mathrm{kpc}}^{2}\,{\mathrm{Myr}}^{-2}$.

In Figure 6, we depict the last 5 Gyr of Pegasus IV's orbit (in various projections) assuming the velocity, distance, and proper motions reported in Table 1 as initial conditions. Notably, the model predicted that Pegasus IV passed its apocenter within the last ∼200 Myr and experienced its last pericentric passage ∼1 Gyr ago.

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To contextualize Pegasus IV's proximity to its orbital apocenter, we computed the ratio: f = (d_GC−r_peri)/(r_apo−r_peri) following Fritz et al. (2018). This ratio quantifies a satellite's proximity to its pericenter (f = 0) or apocenter (f = 1). Assuming Pegasus IV's distance to the Galactic center is d_GC = 89 kpc and adopting the apocenter/pericenter distances given above, we found f = 0.92. This value for the ratio f places Pegasus IV in a regime that is underpopulated compared to the predictions from simple orbital dynamics (e.g., see Figure 5 of Li et al. 2022). Our discovery of Pegasus IV in a previously surveyed region of the sky may support the hypothesis that the dearth of known Milky Way satellite galaxies observed near their apocenters (f ∼ 1) is an observational selection effect (e.g., Fritz et al. 2018; Simon 2018; Li et al. 2022, although see Pace et al. 2022 for a more detailed discussion of this hypothesis).

A number of recently discovered ultra-faint dwarf galaxies have been proposed to be associated with the Large Magellanic Cloud (LMC; e.g., Koposov et al. 2015a; Drlica-Wagner et al. 2015; Erkal & Belokurov 2020; Patel et al. 2020; Correa Magnus & Vasiliev 2022). To assess whether Pegasus IV is a satellite of the LMC, we rewound the system in the combined presence of the LMC and Milky Way potential using the technique described in Erkal & Belokurov (2020). For the Milky Way potential, we used the potential fits of McMillan (2017). We note that we did not select the highest likelihood potential but instead sampled the Milky Way from the posterior chains of McMillan (2017) to account for uncertainties in the potential. We modeled the LMC as a Hernquist profile (Hernquist 1990) with a mass of 1.38 × 10¹¹ M_⊙ and a scale radius of 16.08 kpc, motivated by the results of Erkal et al. (2019). In these simulations, we treated the LMC and Milky Way as particles sourcing their respective potentials and thus account for the reflex motion of the Milky Way in response to the LMC (e.g., Gómez et al. 2015). We modeled the dynamical friction of the Milky Way on the LMC using the approximations in Jethwa et al. 2016. For the LMC's present-day proper motions, distance, and radial velocity we used values provided by Kallivayalil et al. (2013), Pietrzyński et al. 2019, and van der Marel et al. (2002), respectively.

In order to account for uncertainties, we Monte Carlo sampled the present-day observables of Pegasus IV, the Milky Way potential, and the LMC's present-day observables 10,000 times and rewound the satellite for 5 Gyr. ⁴⁰ We computed the energy of Pegasus IV relative to the LMC 5 Gyr ago (as in Erkal & Belokurov 2020), and found that Pegasus IV has a 0.07% chance of having originally been energetically bound to the LMC, suggesting that it is not an LMC satellite. We also considered the approach of Patel et al. (2020) and determined the closest passage of Pegasus IV to the LMC and compared their relative speed to the escape speed of the LMC. With this approach, we found that Pegasus IV passes the LMC at 61 ± 12 kpc, with a relative speed of 363 ± 19 km s⁻¹. This is ∼3 times the escape speed of the LMC, which also suggests that Pegasus IV is not an LMC satellite.

A substantial fraction of the known Milky Way satellite galaxies lies on a thin, corotating plane nearly perpendicular to the Milky Way's stellar disk dubbed the Vast Polar Structure (VPOS; Pawlowski et al. 2012, 2015; Fritz et al. 2018; Li et al. 2021). Adopting the same VPOS parameters as Li et al. (2021), namely, the assumed normal (l_MW, b_MW) = (1693, −2. 8) and angular tolerance θ_inVPOS = 3687, we found it unlikely that Pegasus IV is a VPOS member. The observed angle between the VPOS and the satellite's orbital pole is ${\theta }_{\mathrm{VPOS}}={52.3}_{-19\buildrel{\circ}\over{.} 5}^{+19.8}$ and the probability that the orbital pole lies within θ_inVPOS of the VPOS normal is ∼20%. While this does not rule out the possibility that Pegasus IV is a VPOS member, the currently available phase space measurements do not favor this scenario.

Lastly, we considered whether Pegasus IV might be associated with debris from the Sagittarius dwarf spheroidal galaxy (Sgr dSph) and its extended stellar stream. Considering the Sgr dSph model and associated coordinate system from Law & Majewski (2010), we found that Pegasus IV is located at an angle of β = −529 from the Sgr dSph debris plane. We found a comparably large separation when considering the newer Sgr dSph model from Vasiliev et al. (2021), who additionally incorporated the impact of the LMC when modeling the Sgr dSph's debris stream. We therefore conclude that Pegasus IV is unlikely to be associated with the Sgr dSph.

The Milky Way dwarf spheroidal satellite galaxies are excellent targets for searches for dark matter annihilation or decay products due to their close proximity, astrophysical backgrounds, and large mass-to-light ratios (e.g., Ackermann et al. 2015). The astrophysical component of the dark matter flux from annihilation (decay) is known as the J-factor (D-factor) and depends on the squared (linear) dark matter density along the line of sight. Our framework to calculate J-factors and D-factors follows Pace & Strigari (2019) and is similar to other previous analyses of dSph galaxies (e.g., Bonnivard et al. 2015; Geringer-Sameth et al. 2015). Briefly, we solved for the velocity dispersion in the spherical Jeans equations and compared it to the velocity dispersion from the spectroscopic members to determine the dark matter density profile. We assumed the dark-matter-dominated mass follows an NFW profile, while the stellar distribution follows a Plummer profile. We assumed that stellar anisotropy is constant with radius. We used the results derived in Section 3 for the distance, structural parameters (a_h , ), and associated uncertainties, which were transformed into Gaussian priors. For more details, see Pace & Strigari (2019).

We applied this methodology to the same seven-star (nonvariable) member sample used for our dynamical analysis in Section 4.5. We calculated integrated J-factors of ${\mathrm{log}}_{10}J=17.7\pm 0.8,17.8\pm 0.8,17.9\pm 0.8$ for solid angles of θ = 01, 02, 05 in logarithmic units of GeV² cm⁻⁵. The integrated D-factors are ${\mathrm{log}}_{10}D=16.9\pm 0.4,17.3\,\pm 0.5,17.8\,\pm \,0.6$ for solid angles of θ = 01, 02, and 05 in logarithmic units of GeV cm⁻². The predicted J-factor is ${\mathrm{log}}_{10}J(0\buildrel{\circ}\over{.} 5s)\sim 17.6$ based on velocity dispersion, heliocentric distance, and half-light radius scaling relations and agrees with the full dynamical analysis (Pace & Strigari 2019). This J-factor is not large compared to other ultra-faint dwarfs due primarily to the relatively large distance of Pegasus IV. We note that if Pegasus IV were located at its pericenter (d = 32 kpc), its J-factor would be comparable to the largest J-factors measured for other dwarf galaxies: ${\mathrm{log}}_{10}J(0\buildrel{\circ}\over{.} 5)\sim 18.8$.

RRL-type variable stars are excellent tracers of old, metal-poor stellar populations in the Milky Way halo, and have been identified in nearly every ultra-faint dwarf galaxy (e.g., Greco et al. 2008; Boettcher et al. 2013; Medina et al. 2017; Joo et al. 2018, 2019; Martínez-Vázquez et al. 2019; Vivas et al. 2020; Martínez-Vázquez et al. 2021). According to the empirical relation derived by Martínez-Vázquez et al. (2019), ultra-faint dwarf galaxies with the same absolute magnitude as Pegasus IV (M_V = −4.25) are expected to have between two and four RRLs.

As introduced in Section 2.3, we identified two RRLs in the Gaia and PS1 RRL catalogs within a 2' radius of Pegasus IV's centroid at the time of discovery. The first of these stars (Gaia DR2/EDR3 SOURCE_ID: 1796887082536156928; Gaia G = 20.08 mag) was labeled as an RRab star in both catalogs with a period of 0.7088 days (averaging between the individual catalogs, which agreed at the level of 0.0001 days). This star was identified as a spectroscopic member in Section 4.4 on the basis of its radial velocity. The second of these stars (Gaia DR2/EDR3 SOURCE_ID: 1796890209272433792; G = 20.24 mag) was labeled as an RRc-type variable with a period of 0.31373 days (again averaging between Gaia and PS1, which agreed within 0.00001 days for this star); we do not have a spectrum for this star.

Under the assumption that these stars were bona fide RRL member stars of Pegasus IV, we estimated their absolute magnitudes using the empirical calibration given in Muraveva et al. (2018):

$$\begin{eqnarray}&&{M}_{G}=\left(0.32\pm 0.04\right)[\mathrm{Fe}/{\rm{H}}]+(1.11\pm 0.06).\end{eqnarray} \tag{ 3 }$$

Assuming that the (unknown) RRL metallicities are sampled from the Pegasus IV metallicity distribution function, which we approximate as a Gaussian centered on [Fe/H] = −2.63 dex with variance σ = 0.47 dex, we found that the expected absolute magnitude of the two stars is ${M}_{G}={0.26}_{-0.20}^{+0.18}$ mag, where the uncertainties include contributions from both the sampled RRL metallicity and the errors associated with the coefficients in the Muraveva et al. (2018) relation. From this absolute magnitude, the resulting distance modulus for each of the RRLs was then derived from

$$\begin{eqnarray}&&{(m-M)}_{0}=G-({R}_{G}\times (E(B-V))-{M}_{G},\end{eqnarray} \tag{ 4 }$$

where R_G is the ratio of total-to-selective absorption for the Gaia G filter, which we assumed to be R_G = 2.45 (Wang & Chen 2019). Taking E(B − V) = 0.06 mag for both stars (Table 1), we found ${(m-M)}_{0}\,=\,{19.66}_{-0.18}^{+0.20}$ from the first RRL and for the second, neglecting the errors on G and E(B − V) as they were subdominant to the error on M_G . The average of these distance moduli is (m−M)₀ = 19.74 ± 0.13, in excellent agreement with the distance modulus derived from isochrone fitting, (m−M)₀=19.77 ± 0.03 (stat) ± 0.1 (sys) (Section 3).

The Gaia and PS1 RRL catalogs include an additional RRL located at (α_J2000, δ_J2000) = (328834, 26602), corresponding to a 15.8' separation from Pegasus IV's centroid, or roughly 10 half-light radii (∼0.42 kpc). This star (Gaia DR2/EDR3 SOURCE_ID: 1796879729552126080; Gaia G = 20.12 mag) was flagged in the PS1 catalog as an RRc with a period P = 0.400555 days. Its Gaia EDR3 proper motion (μ_α *, μ_δ ) = (+0.114 ± 0.460, −0.328 ± 0.541) mas yr⁻¹ is consistent with the systemic mean proper motion derived in Section 4.7: (+0.33 ± 0.07, −0.21 ± 0.08) mas yr⁻¹. The distance modulus of this star according to the Muraveva et al. (2018) relation is ${(m-M)}_{0}={19.70}_{-0.18}^{+0.20}$, lying between the distance moduli derived for the other two RRLs discussed in the previous section, and in equally good agreement with the distance modulus derived through isochrone fitting.

These properties suggest that this RRL may be related to Pegasus IV, despite its extreme angular separation. To quantify the possibility that this star is a field RRL, as opposed to a true Pegasus IV member, we integrated the RRL number density radial profile given in Medina et al. (2018) between Galactocentric distances of 80 and 100 kpc. We found that only 0.0075 RRL stars are expected in a 0.25 deg² region around Pegasus IV. Thus, it is very unlikely that this star is a field star, as opposed to a true Pegasus IV member.

RRLs with large angular separations have been observed in the vicinity of several ultra-faint dwarf galaxies (e.g., Vivas et al. 2020; Stringer et al. 2021), and have been proposed to be tidally stripped members of these galaxies. ⁴¹ To assess whether tidal stripping is needed to explain the position of this RRL relative to Pegasus IV, we calculated the system's Jacobi radius following Equation (8.91) of Binney & Tremaine (2008). As explained by Binney & Tremaine (2008), the Jacobi radius approximately corresponds to the expected maximum observed extent of a satellite system in a circular orbit. Adopting the dynamical and structural properties from Table 1, and assuming the simple power-law Milky Way potential from Eadie & Harris (2016), we found that the Jacobi radius for Pegasus IV is ∼0.6 kpc—larger than the projected separation of this RRL from the main body of Pegasus IV (∼0.42 kpc). However, if we instead perform this calculation assuming that Pegasus IV is at its pericenter distance (r_peri = 32 kpc), the Jacobi radius is found to be ∼0.26 kpc, smaller than the observed projected separation. We note, though, that these Jacobi radii are significant underestimates, as they are calculated using the dynamical mass within r_1/2 in absence of a total mass estimate for Pegasus IV.

This latter Jacobi radius estimate admits the possibility that the distant RRL was tidally stripped from the main body of Pegasus IV at a previous pericentric passage, although the close clustering of the confirmed spectroscopic members somewhat disfavors this interpretation. Ultimately, it is difficult to confirm or dispute this star's connection to Pegasus IV without a radial velocity measurement. Wider-area spectroscopic member samples may allow for searches for features suggestive of tidal disruption (e.g., velocity gradients), which would add credence to the tidal origin of this distant star if present. Improved distance estimation for each of the RRLs may also offer further insight into the consistency of this star with the majority of Pegasus IV's members.

Lastly, we note also that there may be yet more RRL members of Pegasus IV, as the Gaia and PS1 RRL catalogs are incomplete at faint magnitudes (e.g., Mateu et al. 2020). Our team has recently obtained deeper Gemini North/GMOS imaging of Pegasus IV (GN-2021B-FT-111; PI: C. Martinez-Vazquez). We therefore defer a more extensive search for RRLs in the central region of Pegasus IV to a future study leveraging these data. These new data will also help disambiguate the nature of the second spectroscopically observed HB star, which appeared to show some signs of variability in the sparsely sampled DELVE data.