##### The Supernova Delay Problem

Simultaneous fits of several line ratios are typically used to determine
$f_{esc}$; this information breaks degeneracies between ionization-bounded and
density-bounded models (Plat et al., 2019) or between density-bounded channels
and density-bounded galaxies (Ramambason et al., 2020). We previously mapped
the gas-phase oxygen abundance and interstellar reddening in J1044+0353 (Peng
et al., 2023). In this paper, we identified a density-bounded ionization cone
based on the anisotropic structure seen an F([O III] $\lambda\lambda
4959,5007)/F(\mbox{[O II]}\lambda\lambda 3726,29)$ flux ratio map. The O32
ratio is high in the starburst region, declines with radius at most position
angles, but remains high throughout a density-bounded ionization cone
extending north-northeast. A key part of this argument relied on [O III]
detections at large separations, 1.5 - 2.5 kpc, from the starburst region
where no [O II] emission was detected, yielding a lower limit on O32. Along
the line of sight to J1044+0353, only a small fraction of the starburst is
covered by low column density H I (Hu et al., 2023). Hence, we argue that the
LyC leakage from J1044+0353 is primarily in the direction of the ionization
cone. This anisotropic escape results from the ionizing luminosity, produced
in the starburst region, encountering the cavity excavated by a bipolar
outflow launched from region FEN. The older stellar age of FEN makes this
region a strong source of mechanical feedback even though its current
production rate of ionizing photons is negligible compared to the compact
starburst. At least in J1044+0353, the star formation history solves the
supernova delay problem.

##### The Catastrophic Cooling Problem

In numerical simulations, galactic winds form through superbubble breakthrough
and blowout (Mac Low & McCray, 1988; Mac Low et al., 1989). Remarkably, we
show that superbubble breakthrough and blowout imprint distinct signatures on
the starburst spectrum. Toward the starburst region, the [O III] and H$\beta$
line profiles show a broad (213 ${\rm km~{}s}^{-1}$ FWHM) and a very broad
(752 ${\rm km~{}s}^{-1}$ FWHM) components. We mapped these components
individually. The broad component has the higher luminosity, is spatially
extended, and is centered on the starburst region. The two most prominent
superbubbles within this region extend 350 and 400 pc, respectively, from the
compact star clusters labeled Region 304 and 303. Growing these bubbles
requires shell velocities $v_{SW}\approx 100(t/2{\rm~{}Myr})^{-1}$ ${\rm
km~{}s}^{-1}$ and $v_{NW}\approx 120(t/2{\rm~{}Myr})^{-1}$ ${\rm
km~{}s}^{-1}$. The starburst age from Olivier et al. (2022) therefore produces
a good match between the shell velocity and the blueshift of UV absorption
lines (Xu et al., 2022). We therefore suggest that the outflowing gas
identified by Xu et al. (2022) resides in a superbubble shell rather than a
multiphase galactic wind; the actual bipolar wind form J1044+0353 is outside
their spectroscopic aperture. The superbubble dynamics constrain the supernova
rate in star clusters believed to have formed some very massive stars, a
population where empirical constraints on supernova rates are needed (Fryer et
al., 2023). A filamentary arm extending 1 kpc to the west of the starburst
appears only in the lower intensity, very-broad component. Close examination
of the Hubble  H$\alpha$ image reveals faint H$\alpha$ emission on each side
of this very-high velocity gas. These features are consistent with superbubble
blowout, and the velocity of the very-broad component matches expectations for
the velocity of a free-flowing, hot wind (Chevalier & Clegg, 1985). These
observations suggest a picture where the young starburst is starting to launch
a galactic wind. A comparison to the Thompson et al. (2016) wind model
demonstrated that the radiative cooling from this nascent wind is not
catastrophic. The [O III] luminosity of the wind is higher than single-phase
wind models can explain, however; and the discrepancy may provide a useful
target for multi-phase wind models. These results aid the interpretation of
the emission-line wings detected in 25-40% of redshift $3<z<9$ galaxies
(Carniani et al., 2023; Xu et al., 2023).

The galaxy population believed to drive cosmic reionization had low masses,
very low abundances of heavy elements, and high star-formation efficiencies
(Ouchi et al., 2009; Wise & Cen, 2009; Bouwens et al., 2015b). JWST is finding
a large population of such galaxies via their strong [O III] emission (Oesch
et al., 2023; Tang et al., 2023; Matthee et al., 2023). In this paper, we have
demonstrated that the star formation history across a local analog,
J1044+0353, plays a critical role in forming pathways for LyC escape. The LyC
photons generated by the very young starburst encounter an ISM whose structure
has already been modified by a galactic wind launched 10-20 Myr prior from a
post-starburst region. These results support the idea that galaxy orientation
determines which strong W([O III]) have directly detected LyC (Tang et al.,
2021). We suggest that due to anisotropic LyC escape, spatial mapping of the
O32 ratio may turn out to be a more accurate indicator of LyC escape than
direct detection of the Lyman cotinuum, which requires an optically thin
channel to be aligned with our sightline.

Our schematic picture of J1044+0353 is qualitatively similar to the scenario
seen in full-physics simulations of reionization-era galaxies; the SFR is
highly variable in time, and LyC escape is anisotropic (Martin-Alvarez et al.,
2023; Witten et al., 2024). Bursty star formation histories at redshift $z\sim
10-13$ also offer an explanation (Muñoz et al., 2023; Sun et al., 2023b, a)
for the surprising luminosity excess observed (Naidu et al., 2022; Finkelstein
et al., 2023; Casey et al., 2023). The rapid build up of galaxies during the
reionization era produces many galaxy – galaxy mergers (Witten et al., 2024).
In contrast, it is not obvious that local analogs selected to resemble high-
redshift galaxies based on their spectral properties would share this
environmental property. Indeed, the stellar mass of J1044+0353 is well below
the limits of pair surveys (Stierwalt et al., 2017; Besla et al., 2018), so
the probability of a dwarf – dwarf merger in the $M_{*}\sim 10^{6}$
${\rm~{}M_{\odot}}$ to $10^{7}$ ${\rm~{}M_{\odot}}$ range is not constrained
observationally. J1044+0353 has neither a luminous companion nor any known
tidal tails, i.e. features that would confirm an interaction or merger
hypothesis. We note, however, the unexplained presence of a band of high
extinction, elongated perpendicular to a line drawn from SB to FEN (Peng et
al., 2023); the coincident string of older star clusters there resembles Mrk
709 (Kimbro et al., 2021). Hence we suggest that a dwarf – dwarf merger
triggered the recent star formation activity in J1044+0353. Mergers of dwarf
galaxies have been observed to produce more widely distributed star formation
than mergers of massive galaxies, which funnel gas into a central starburst
(Privon et al., 2017), and we would expect the first passage to produce a
starburst in both dwarfs (Kimbro et al., 2021). A merger scenario predicts a
positive radial age gradient in the starburst region, similar to the one
measured in 30 Doradus (Brandl et al., 1996; Cignoni et al., 2015) and some
interstellar gas flung out to large distances (Pearson et al., 2016, 2018) and
possibly detectable via 21-cm mapping.

We thank George Privon, Zirui Chen, Mordecai Mac-Low, Drummond Fielding, and
Norman Murray for discussions that improved the paper. Pierre Thibodeaux
developed the adapative binning code used here. The data presented herein were
obtained at the W. M. Keck Observatory, which is operated as a scientific
partnership among the California Institute of Technology, the University of
California and the National Aeronautics and Space Administration. The
Observatory was made possible by the generous financial support of the W. M.
Keck Foundation. The authors wish to recognize and acknowledge the very
significant cultural role and reverence that the summit of Maunakea has always
had within the indigenous Hawaiian community. We are most fortunate to have
the opportunity to conduct observations from this mountain. This work was
completed at the Aspen Center for Physics, which is supported by National
Science Foundation grant PHY-2210452. This work was supported by the National
Science Foundation through AST-1817125.

Table 1: Basic Properties of J0144+0353 Property | Derived | Unit | Reference
---|---|---|---
W([O III]) | $1040\pm 20$ | Å | NSA Catalog
M${}_{R}-5\log h$ | -14.95 | mag | NSA CatalogaaScaled to the distance adopted in this paper gives M${}_{FUV}=-16.23$ M${}_{R}=-15.78$.
M${}_{FUV}-5\log h$ | -15.40 | mag | NSA CatalogaaScaled to the distance adopted in this paper gives M${}_{FUV}=-16.23$ M${}_{R}=-15.78$.
E(B-V) (Galactic) | 0.0366 | mag | Schlafly & Finkbeiner (2011a)
AV | $0.20^{+0.10}_{-0.10}$ | mag | Peng et al. (2023)
SFRTot (SED) | $0.26^{+0.07}_{-0.07}$ |  ${\rm~{}M_{\odot}}$ yr-1 |  Berg et al. (2022)bbAssumes a luminosity distance of 55 Mpc. Scale by 1.15 to match the distance adopted throughout this paper.
SFRAp (SED) | $0.13^{+0.02}_{-0.03}$ |  ${\rm~{}M_{\odot}}$ yr-1 |  Berg et al. (2022)bbAssumes a luminosity distance of 55 Mpc. Scale by 1.15 to match the distance adopted throughout this paper.
$\log{\rm M}_{*,Tot}$ (SED) | $6.8^{+0.41}_{-0.26}$ | ${\rm~{}M_{\odot}}$ |  Berg et al. (2022)bbAssumes a luminosity distance of 55 Mpc. Scale by 1.15 to match the distance adopted throughout this paper.
$\log{\rm M}_{*,Ap}$ (SED) | $6.13^{+0.23}_{-0.09}$ | ${\rm~{}M_{\odot}}$ |  Berg et al. (2022)bbAssumes a luminosity distance of 55 Mpc. Scale by 1.15 to match the distance adopted throughout this paper.
$r_{50}(UV)$ | 0.38 | arcseconds | Xu et al. (2022)
SFR/Area | $0.6^{+0.17}_{-0.14}$ |  ${\rm~{}M_{\odot}}$ yr-1 kpc-2 | Xu et al. (2022)
log Specific SFR | $-7.39^{+0.43}_{-0.28}$ | yr-1 | Xu et al. (2022)
$v_{circ}$ | $15.6^{+3.0}_{-3.7}$ | ${\rm km~{}s}^{-1}$ | Xu et al. (2022)
$v_{out}$ (UV, LOS) | $52\pm 13$ | ${\rm km~{}s}^{-1}$ | Xu et al. (2022)
$v_{out,fwhm}$ (UV) | $123\pm 24$ | ${\rm km~{}s}^{-1}$ | Xu et al. (2022)
vout ([O III]) | +50 to -40 | ${\rm km~{}s}^{-1}$ | Peng et al. (2023)
$T_{e}$ ([O III]) | $19,200\pm 200$ | K | Berg et al. (2021)
$T_{e}$ ([O II]) | $19,100\pm 1500$ | K | Berg et al. (2021)
12 + $\log$ (O/H) | $7.45\pm 0.03$ |  | Berg et al. (2022)
$n_{e}$ ([S II]) | $200\pm 40$ | cm-3 | Berg et al. (2022)
$n_{e}$ (C III]) | $<9000$ | cm-3 | Berg et al. (2021)
$\log\mbox{$\xi_{ion}$}$ | $25.8\pm 0.9$ | Hz erg-1 | Olivier et al. (2022)
SB Age (UV Spectrum, BPASS) | $4.03\pm 0.85$ | Myr | Olivier et al. (2022)
SB Age (UV Spectrum, SB 99) | $1.26\pm 2.36$ | Myr | Olivier et al. (2022)
SB Age (Balmer EW) | $6.480\pm 0.005$ | Myr | Peng et al. (2023)
FEN Age (Balmer EW) | $7.30\pm 0.05$ | Myr | Peng et al. (2023)
NEN/MEN Age (Balmer EW) | $7.2\pm 0.1$ | Myr | Peng et al. (2023)
Table 2: Starburst Photometry and SED Fitting

| SB | SW-shell | NW-shell | 301 | 302 | 303 | 304
---|---|---|---|---|---|---|---
F125LP (mJy) | $0.1181$ | $0.0079$ | $0.00506$ | 0.00324 | 0.00324 | 0.0312 | 0.0359
F150LP (mJy) | $0.1362$ | $0.0089$ | $0.0061$ | 0.00334 | 0.00334 | 0.0379 | 0.0436
F165LP (mJy) | $0.1334$ | $0.0100$ | $0.0065$ | 0.00333 | 0.00333 | 0.0362 | 0.0397
F336W (mJy) | $0.1729$ | $0.0109$ | $0.0064$ | 0.00635 | 0.00635 | 0.0532 | 0.055
F438W (mJy) | $0.174$ | $0.0119$ | $0.0073$ | 0.00636 | 0.00636 | 0.0521 | 0.0527
F606W (mJy) | $0.322$ | $0.0173$ | $0.0126$ | 0.0132 | 0.0132 | 0.101415 | 0.0999
F665N (mJy) | $1.5615$ | $0.077$ | $0.0631$ | 0.067 | 0.067 | 0.502318 | 0.48167
F814W (mJy) | $0.1105$ | $0.0086$ | $0.0056$ | 0.00408 | 0.00408 | 0.0314 | 0.0308
$M$ ($10^{5}$ ${\rm~{}M_{\odot}}$)aa Beagle fits using single-age stellar population and smc reddening. Age defined as the SSP age.  | $11.78^{+0.34}_{-0.20}$ | $1.31^{+0.29}_{-0.19}$ | $0.41^{+0.01}_{-0.01}$ | $0.52^{+0.04}_{-0.02}$ | $1.04^{+0.06}_{-0.05}$ | $3.67^{+0.18}_{-0.17}$ | $3.44^{+0.15}_{-0.15}$
$t_{ssp}$ (Myr)aa Beagle fits using single-age stellar population and smc reddening. Age defined as the SSP age.  | $2.98^{+0.01}_{-0.01}$ | $5.02^{+0.33}_{-0.22}$ | $3.03^{+0.01}_{-0.01}$ | $2.95^{+0.02}_{-0.03}$ | $2.97^{+0.02}_{-0.03}$ | $2.97^{+0.02}_{-0.02}$ | $2.98^{+0.01}_{-0.02}$
$\tau_{e}ff$aa Beagle fits using single-age stellar population and smc reddening. Age defined as the SSP age.  | $0.228^{+0.007}_{-0.007}$ | $0.248^{+0.032}_{-0.026}$ | $0.181^{+0.004}_{-0.007}$ | $0.337^{+0.012}_{-0.011}$ | $0.249^{+0.010}_{-0.011}$ | $0.254^{+0.011}_{-0.011}$ | $0.213^{+0.010}_{-0.010}$
$M$ ($10^{5}$ ${\rm~{}M_{\odot}}$)bb Beagle fits using single-age stellar population and CF00 reddening.  | $11.43^{+0.62}_{-0.50}$ | $1.22^{+0.15}_{-0.16}$ | $0.42^{+0.02}_{-0.02}$ | $0.54^{+0.03}_{-0.03}$ | $1.07^{+0.06}_{-0.05}$ | $3.86^{+0.20}_{-0.20}$ | $3.52^{+0.15}_{-0.15}$
$t_{ssp}$ (Myr)bb Beagle fits using single-age stellar population and CF00 reddening.  | $2.99^{+0.01}_{-0.02}$ | $4.88^{+0.18}_{-0.17}$ | $3.03^{+0.01}_{-0.01}$ | $2.95^{+0.02}_{-0.03}$ | $2.96^{+0.02}_{-0.03}$ | $2.96^{+0.02}_{-0.03}$ | $2.98^{+0.01}_{-0.02}$
$\tau_{e}ff$bb Beagle fits using single-age stellar population and CF00 reddening.  | $0.159^{+0.012}_{-0.015}$ | $0.240^{+0.019}_{-0.025}$ | $0.181^{+0.017}_{-0.013}$ | $0.348^{+0.012}_{-0.012}$ | $0.252^{+0.011}_{-0.012}$ | $0.261^{+0.012}_{-0.011}$ | $0.217^{+0.010}_{-0.011}$
$M$ ($10^{5}$ ${\rm~{}M_{\odot}}$)cc Beagle fits using const SFH and smc reddening. Age defined by the maximum stellar age.  | $11.74^{+0.63}_{-0.43}$ | $2.09^{+0.29}_{-0.15}$ | $0.78^{+0.19}_{-0.19}$ | $0.56^{+0.03}_{-0.03}$ | $1.12^{+0.05}_{-0.05}$ | $4.07^{+0.16}_{-0.17}$ | $3.80^{+0.15}_{-0.17}$
$t_{max}$ (Myr)cc Beagle fits using const SFH and smc reddening. Age defined by the maximum stellar age.  | $3.37^{+0.06}_{-0.08}$ | $21.88^{+3.83}_{-4.73}$ | $10.23^{+2.95}_{-2.98}$ | $3.24^{+0.05}_{-0.08}$ | $3.24^{+0.05}_{-0.05}$ | $3.24^{+0.05}_{-0.05}$ | $3.24^{+0.05}_{-0.05}$
$\tau_{e}ff$cc Beagle fits using const SFH and smc reddening. Age defined by the maximum stellar age.  | $0.156^{+0.011}_{-0.014}$ | $0.220^{+0.012}_{-0.012}$ | $0.198^{+0.023}_{-0.025}$ | $0.290^{+0.013}_{-0.013}$ | $0.201^{+0.012}_{-0.012}$ | $0.208^{+0.012}_{-0.011}$ | $0.164^{+0.012}_{-0.013}$
$M$ ($10^{5}$ ${\rm~{}M_{\odot}}$)dd Beagle fits using const SFH and CF00 reddening.  | $12.01^{+0.75}_{-0.53}$ | $1.94^{+0.32}_{-0.29}$ | $0.74^{+0.18}_{-0.14}$ | $0.58^{+0.03}_{-0.03}$ | $1.15^{+0.05}_{-0.05}$ | $4.13^{+0.16}_{-0.17}$ | $3.86^{+0.18}_{-0.17}$
$t_{max}$ (Myr)dd Beagle fits using const SFH and CF00 reddening.  | $3.39^{+0.06}_{-0.11}$ | $17.34^{+2.80}_{-2.75}$ | $9.40^{+2.32}_{-2.74}$ | $3.23^{+0.07}_{-0.07}$ | $3.24^{+0.05}_{-0.05}$ | $3.24^{+0.05}_{-0.05}$ | $3.26^{+0.05}_{-0.05}$
$\tau_{e}ff$dd Beagle fits using const SFH and CF00 reddening.  | $0.159^{+0.012}_{-0.016}$ | $0.318^{+0.052}_{-0.047}$ | $0.195^{+0.040}_{-0.025}$ | $0.294^{+0.015}_{-0.013}$ | $0.204^{+0.013}_{-0.012}$ | $0.211^{+0.012}_{-0.011}$ | $0.168^{+0.013}_{-0.012}$

eefootnotetext:  The photometric errors are dominated by systematic errors
which are estimated to be 5%.

fffootnotetext: Comparison of the parameters for the four fitted models
illustrate their systematic uncertainties. The ionization parameter $\log(U)$,
stellar metallicity $\log(Z/Z_{\odot})$, and gas-phase metallicity
$\log(Z_{gas}/Z_{\odot})$ were fixed at -1.75, -1.296, and -1.296. A Chabrier
IMF from 0.1 to 100 ${\rm~{}M_{\odot}}$ was adopted. Departures from these
values would increase the systemic error. The statistical errors given were
obtained by the 16th and 84th percentile of the posterior distribution of the
parameters in comparison with the median (50th percentile) value.

Table 3: KCWI Measurements of Radiative Feedback

Measurementb,cb,cfootnotemark:  | Unit | J104457+035313
---|---|---
|  RC(H+) ApertureaaThe symbols RC(H+), RC(He+2), and RC(He+) denote the H+, He+2, and He+ Strömgren radius, respectively. The symbol RX denotes the extent of the filamentary emission; see Fig. 6. |
RC(H+) | pc | 1990
RX(H+) | pc | $>3340$
F(H$\beta$) | ergs s-1 cm-2 | $4.600\pm 0.002\times 10^{-14}$
F([O II] $\lambda 3726$) | ergs s-1 cm-2 | $93.80\pm xx\times 10^{-16}$
F([O II] $\lambda 3729$) | ergs s-1 cm-2 | $122.49\pm xx\times 10^{-16}$
L’(H$\beta$) | ergs s-1 | $2.54\pm 0.15\times 10^{40}$
Q(H) | s-1 | $5.51\pm 0.22\times 10^{52}$
$\langle n_{e}^{2}\rangle^{1/2}$ | ${\rm cm^{-3}}$ | 0.69
$n_{e}(\mbox{[O II]})$ | ${\rm cm}^{-3}$ | $161^{+7}_{-10}$
$\epsilon$([O II]) | $\cdots$ | $1.8\times 10^{-5}$
M${}_{FUV}^{Tot}$ | AB Mag | $-16.68\pm 0.17$
$\log\mbox{$\xi_{ion}({\rm H})$}_{Tot}$ | Hz erg-1 | $25.43^{+0.10}_{-0.13}$
$\log\mbox{$\xi_{ion}$}(\mbox{He${}^{+}$})_{Tot}$ | Hz erg-1 | $23.39^{+0.11}_{-0.14}$
|  RC(He+2) ApertureaaThe symbols RC(H+), RC(He+2), and RC(He+) denote the H+, He+2, and He+ Strömgren radius, respectively. The symbol RX denotes the extent of the filamentary emission; see Fig. 6. |
RC(He+2) | pc | 610
RX(He+2) | pc | 730
W([O III]) | (Å) | 1230
F(H$\beta$) | ergs s-1 cm-2 | $3.752\pm 0.002\times 10^{-14}$
F(He II $\lambda 4686$) | ergs s-1 cm-2 | $7.24\pm 0.03\times 10^{-16}$
F([O II]$\lambda 3726$) | ergs s-1 cm-2 | $76.08\pm xx\times 10^{-16}$
F([O II]$\lambda 3729$) | ergs s-1 cm-2 | $97.68\pm xx\times 10^{-16}$
F([O III]$\lambda 4959$) | ergs s-1 cm-2 | $600.6\pm 0.1\times 10^{-16}$
F([Ar IV] $\lambda 4740$) | ergs s-1 cm-2 | $6.47\pm xx\times 10^{-16}$
F([Ar IV] $\lambda 4711$) | ergs s-1 cm-2 | $7.75\pm xx\times 10^{-16}$
L’(He II $\lambda 4686$) | ergs s-1 | $4.0\pm 0.2\times 10^{38}$
Q(He+) | s-1 | $4.89\pm 0.47\times 10^{50}$
$\langle n_{e}^{2}\rangle^{1/2}$(He II) | ${\rm cm^{-3}}$ | 0.50
$\langle n_{e}^{2}\rangle^{1/2}$(H I) | ${\rm cm^{-3}}$ | 3.7
$n_{e}(\mbox{[O II]})$ | ${\rm cm}^{-3}$ | $185^{+9}_{-11}$
$\epsilon$([O II]) | $\cdots$ | $4\times 10^{-4}$
$n_{e}(\mbox{[Ar IV]}))$ | ${\rm cm}^{-3}$ | $1440^{+44}_{-128}$
$\epsilon_{h}$([Ar IV]) | $\cdots$ | $1.2\times 10^{-7}$
M${}_{FUV}^{SB}$ | AB Mag | $-16.32\pm 0.17$
$\log\mbox{$\xi_{ion}({\rm H})$}_{SB}$ | Hz erg-1 | $25.58^{+0.098}_{-0.13}$
$\log\mbox{$\xi_{ion}$}(\mbox{He${}^{+}$})_{SB}$ | Hz erg-1 | $23.53^{+0.11}_{-0.14}$
|  RC(He+) ApertureaaThe symbols RC(H+), RC(He+2), and RC(He+) denote the H+, He+2, and He+ Strömgren radius, respectively. The symbol RX denotes the extent of the filamentary emission; see Fig. 6. |
RC(He+) | pc | 650
RX(He+) | pc | 1420
F(He I $\lambda 4471$) | ergs s-1 cm-2 | $1.456\pm 0.004\times 10^{-15}$
L’(He I $\lambda 4471$) | ergs s-1 | $8.35\pm 0.80\times 10^{38}$

bbfootnotetext: Line fluxes given before extinction correction. Line
luminosities, L’, have been corrected for extinction. Ionizing photon
luminosities $Q(\mbox{He${}^{+}$})$ and $Q(H)$ computed at $h\nu>54.4$ eV and
$h\nu>13.6$ eV, respectively. Recombination rates are computed at at the
temperature measured in the [O III] zone, $T_{e}\approx 2.0\times 10^{4}$ K.
The ionizing photon efficiency, $\xi_{ion}$, is computed from Equation 6.1,
using the absolute UV magnitude at 1500 Å.

ccfootnotetext: Electron density, $n_{e}$, computed from the doublet ratio of
the ion indicated; spectral apertures are identified by name. Root-mean-square
electron density, $\langle n_{e}^{2}\rangle^{1/2}$: computed from RC(H+) and
H$\beta$ luminosity within RC(H+), computed within RC(He+2) for the ionization
zones probed by He II $\lambda 4686$ as well as H$\beta$.

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