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The rapid growth which the field has experienced in the past (wo decades was to a very large extent stimulated by (he idea that high energv nuclear collisions will produce droplets of stvonegly interacting matter - droplets large enough and long-lived enough to allow a study of the predictions which QCD makes lor macroscopic svstenis.
The rapid growth which the field has experienced in the past two decades was to a very large extent stimulated by the idea that high energy nuclear collisions will produce droplets of strongly interacting matter - droplets large enough and long-lived enough to allow a study of the predictions which QCD makes for macroscopic systems.
Moreover. it is expected that the conditions provided in (hese interactions will sullice for quark plasma formation.
Moreover, it is expected that the conditions provided in these interactions will suffice for quark plasma formation.
Hence the study of strongly interacting matter has today a experimental side: (his. in turn. has stimulated much of the subsequent theoretical development.
Hence the study of strongly interacting matter has today a multi-faceted experimental side; this, in turn, has stimulated much of the subsequent theoretical development.
The relevant experiments were inilially denoted as ullra-relativistic nucleus-nucleus collisions: thev are often. not «quite correctly. also called heavy ion collisions (au ion filly stripped of its electrons is a nucleus).
The relevant experiments were initially denoted as ultra-relativistic nucleus-nucleus collisions; they are often, not quite correctly, also called heavy ion collisions (an ion fully stripped of its electrons is a nucleus).
The studies began al Brookhaven National Laboratory (BNL) near New York and at the European Center for Nuclear Research (CERN) near Geneva around 1986/87.
The studies began at Brookhaven National Laboratory (BNL) near New York and at the European Center for Nuclear Research (CERN) near Geneva around 1986/87.
The first collisions had light nuclei (oxveen. silicon. sulphur) hitting heavy. targets (gold. uranium). since light ions could be dealt with using injectors already. existing at BNL and CERN.
The first collisions had light nuclei (oxygen, silicon, sulphur) hitting heavy targets (gold, uranium), since light ions could be dealt with using injectors already existing at BNL and CERN.
The successful analvsis of these experiments provided the basis and motivation for the construction of new injectors of truly heavy nuclei. gold at BNL and lead at CERN: they came into operation in (he middle 1990s.
The successful analysis of these experiments provided the basis and motivation for the construction of new injectors of truly heavy nuclei, gold at BNL and lead at CERN; they came into operation in the middle 1990's.
These early fixed target experiments were carried oul al a center of mass enerev of around 5 GeV per nucleon-nucleon collision at the BNL-AGS and around 20 GeV at the CERN-SPS.
These early fixed target experiments were carried out at a center of mass energy of around 5 GeV per nucleon-nucleon collision at the BNL-AGS and around 20 GeV at the CERN-SPS.
At the turn of millenium. the first dedicated nuclear accelerator. the Relativistie Leavy Ion Collider RING. started taking data. al BNL. with a center of mass energy a factor ten higher. al 200 GeV per nucleon-nucleon collision.
At the turn of millenium, the first dedicated nuclear accelerator, the Relativistic Heavy Ion Collider RHIC, started taking data at BNL, with a center of mass energy a factor ten higher, at 200 GeV per nucleon-nucleon collision.
And coming soon now. the Large Lacon Collider LUC at CERN will bring the center-of-mass energy. lor nuclear collisions up to 5500 GeV. or 5.5 ΤΟΝ. The work of the different experimental groups working at these facilies has provided an immense wealth of data. ancl Chere is little doubt today that in such collisions comparatively large svstems of higher energy density are formed than have ever been studied in thie laboratory before.
And coming soon now, the Large Hadron Collider LHC at CERN will bring the center-of-mass energy for nuclear collisions up to 5500 GeV, or 5.5 TeV. The work of the different experimental groups working at these facilities has provided an immense wealth of data, and there is little doubt today that in such collisions comparatively large systems of higher energy density are formed than have ever been studied in the laboratory before.
The detailed analvses of the results have also shown. however. that a number οἱ
The detailed analyses of the results have also shown, however, that a number of
surrounding medium.
surrounding medium.
Phe voungest star forming regions are ound mostly to the cast west of NGC 1512. while regions owards the north south of NGC 1512 are generally older.
The youngest star forming regions are found mostly to the east west of NGC 1512, while regions towards the north south of NGC 1512 are generally older.
‘This east-west (voung) versus north-south (old) svmmetry might indicate the passage of NGC 1510 as it is accreted w NGC 1512.
This east-west (young) versus north-south (old) symmetry might indicate the passage of NGC 1510 as it is accreted by NGC 1512.
This interaction might have (1) triggered the xw in NGC 1512 (unless this was the result. of previous interactions or minor mergers) causing gas within the co-rotation radius to Low towards the nuclear region. thus ooviding fuel for continuous star formation. (2) allected he spiral arm pattern causing broadening and splitting as well as enhanced star formation. and (3) led to the ejection of material to large radii where it may become unbound. orming dense clumps able to form new stars.
This interaction might have (1) triggered the bar in NGC 1512 (unless this was the result of previous interactions or minor mergers) causing gas within the co-rotation radius to flow towards the nuclear region, thus providing fuel for continuous star formation, (2) affected the spiral arm pattern causing broadening and splitting as well as enhanced star formation, and (3) led to the ejection of material to large radii where it may become unbound, forming dense clumps able to form new stars.
Evidence of he latter is the observation of twogalarics at he outermost regions of the NGC 1512/1510 system.
Evidence of the latter is the observation of two at the outermost regions of the NGC 1512/1510 system.
‘Phere ave numerous ways to estimate the star formation rate (SER) of a galaxy.
There are numerous ways to estimate the star formation rate (SFR) of a galaxy.
Vo study the global and local SERS. we use a range of line and continuum measurements at cillcrent wavelengths. (ultraviolet. optical. infrared. ancl racio)
To study the global and local SFRs, we use a range of line and continuum measurements at different wavelengths (ultraviolet, optical, infrared, and radio).
A combination of these data together with an understanding
A combination of these data together with an understanding
to R—0; and at much larger 7 the relative amplitude can be damped at smaller R making orbit crossings more concentrated in R.
to $R\rightarrow0$; and at much larger $\eta$ the relative amplitude can be damped at smaller $R$ making orbit crossings more concentrated in $R$.
In Figure 7,, we compare the mode structure (at otherwise similar properties) to the gravitational softening or disk thickness.
In Figure \ref{fig:m1.4}, we compare the mode structure (at otherwise similar properties) to the gravitational softening or disk thickness.
Atfixed w and other disk properties, the modes in colder disks are higher-k (more tightly wound).
At $\omega$ and other disk properties, the modes in colder disks are $k$ (more tightly wound).
This is also evident in the simulations in ? (see their Figures 2 44).
This is also evident in the simulations in \citet{hopkins:zoom.sims} (see their Figures 2 4).
This is because, for almost all the modes of interest, the mode has at least some contributions from the short-branch regime — they are analogous to the “p-modes” in ?, both in that the pressure/softening effects are non-negligible, in that the characteristic modes are trailing (the ΚΚ>0 branch of the p-modes remains ΚΚ>0 after refraction), and in that they have positive (prograde) pattern speeds.
This is because, for almost all the modes of interest, the mode has at least some contributions from the short-branch regime – they are analogous to the “p-modes” in \citet{tremaine:slow.keplerian.modes}, both in that the pressure/softening effects are non-negligible, in that the characteristic modes are trailing (the $kR>0$ branch of the p-modes remains $kR>0$ after refraction), and in that they have positive (prograde) pattern speeds.
Because these modes depend on some non-zero 3, as the modes themselves heatup the disk when they go non-linear (ultimately stabilizing it at some Q threshold), they can become more global.
Because these modes depend on some non-zero $\beta$, as the modes themselves heatup the disk when they go non-linear (ultimately stabilizing it at some $Q$ threshold), they can become more global.
This has the effect (at fixed mode amplitude) of actually increasing the efficiency of the modes at driving large eccentricity and inflows to small radii, although the mode growth rates are lower.
This has the effect (at fixed mode amplitude) of actually increasing the efficiency of the modes at driving large eccentricity and inflows to small radii, although the mode growth rates are lower.
However, most of this difference depending on disk thickness is concentrated at small softening; once moderate disk thickness >0.05 is reached, the effect of the modes is actually fairly weakly dependent—0.1 on the thickness.
However, most of this difference depending on disk thickness is concentrated at small softening; once moderate disk thickness $\gtrsim0.05-0.1$ is reached, the effect of the modes is actually fairly weakly dependent on the thickness.
Figure 8 compares the mode structure at fixed η and f but varying Mais. (and correspondingly a).
Figure \ref{fig:m1.5} compares the mode structure at fixed $\eta$ and $\beta$ but varying $M_{\rm disk}$ (and correspondingly $a$ ).
At low Mais./Mau~0.1— 0.3, the first unstable modes to appear are, unsurprisingly, the maximally unstable modes with large |kR|&1/8.
At low $M_{\rm disk}/M_{\rm BH}\sim0.1-0.3$ , the first unstable modes to appear are, unsurprisingly, the maximally unstable modes with large $|kR|\approx1/\beta$.
As Mais increases, the spectrum of unstable modes expands to include longer-wavelength modes, and by Maia.2Mau includes very global modes.
As $M_{\rm disk}$ increases, the spectrum of unstable modes expands to include longer-wavelength modes, and by $M_{\rm disk}\gtrsim M_{\rm BH}$ includes very global modes.
Once Mai>>Mpu, the structure at the radii we are interested in — iin the quasi-Keplerian regime ~Ro, is essentially independent of Mais (it is identical to the case of an infinite power-law disk).
Once $M_{\rm disk}\gg M_{\rm BH}$, the structure at the radii we are interested in – in the quasi-Keplerian regime $\sim R_{0}$, is essentially independent of $M_{\rm disk}$ (it is identical to the case of an infinite power-law disk).
Of course, there will in this limit be other modes at larger radii that are simply standard disk modes, but we are not interested in these behaviors.
Of course, there will in this limit be other modes at larger radii that are simply standard disk modes, but we are not interested in these behaviors.
Our analysis thus far is restricted to the linear regime.
Our analysis thus far is restricted to the linear regime.
To see whether our key conclusions are robust in the non-linear regime, with gas+stellar systems (albeit still stellar-dominated), in the presence of inflow, star formation, feedback, and a non-trivial potential, we briefly compare to the mode structure in the simulations of self-consistently formed nuclear stellar disks from 9
To see whether our key conclusions are robust in the non-linear regime, with gas+stellar systems (albeit still stellar-dominated), in the presence of inflow, star formation, feedback, and a non-trivial potential, we briefly compare to the mode structure in the simulations of self-consistently formed nuclear stellar disks from \citet{hopkins:m31.disk}.
Figure 9 illustrates this in a typical nuclear scale simulation.
Figure \ref{fig:nuclear.mode.origins} illustrates this in a typical nuclear scale simulation.
We plot the enclosed disk Ma(<R) and BH mass, mode pattern speed 9, (and circular speed (), and mode amplitudes, as a function of time in a system that is initially smooth hhas no m= perturbation).
We plot the enclosed disk $M_{d}(<R)$ and BH mass, mode pattern speed $\Omega_{p}$ (and circular speed $\Omega$ ), and mode amplitudes, as a function of time in a system that is initially smooth has no $m=1$ perturbation).
The mode first appears at some radii ~Rerit, where1 Ma/Msn~1, with Q,~O(Rai).
The mode first appears at some radii $\sim R_{\rm crit}$, where $M_{d}/M_{\rm BH}\sim1$, with $\Omega_{p}\sim\Omega(R_{\rm crit})$.
At early times, the inner disk profile is quite shallow (or even hollow), because no inflow has yet reached the center; we see the resulting cutoff in the range of the mode at small radius.
At early times, the inner disk profile is quite shallow (or even hollow), because no inflow has yet reached the center; we see the resulting cutoff in the range of the mode at small radius.
Two things work to push this range inwards.
Two things work to push this range inwards.
First, the mode slows down at early times, seen in Figure 9..
First, the mode slows down at early times, seen in Figure \ref{fig:nuclear.mode.origins}.
This occurs both via angular momentum exchange with the bulge and disk at somewhat larger (~ 100pc) radii (a resonant process not included in our analysis), and via direct carrying of some of the angular momentum in wavepackets in the gas after reflection off the inner radius above (through the OLR, apparent in the WKB treatment).
This occurs both via angular momentum exchange with the bulge and disk at somewhat larger $\sim100\,$ pc) radii (a resonant process not included in our analysis), and via direct carrying of some of the angular momentum in wavepackets in the gas after reflection off the inner radius above (through the OLR, apparent in the WKB treatment).
This slowdown occurs while the system is in transition from the overstable growth phase to the nonlinear, quasi-steady state.
This slowdown occurs while the system is in transition from the overstable growth phase to the nonlinear, quasi-steady state.
It halts once the €), is significantly below Q(Rait); in these simulations the mode pattern speeds tend to stabilize at values ~1-5kms Ρο.
It halts once the $\Omega_{p}$ is significantly below $\Omega(R_{\rm crit})$; in these simulations the mode pattern speeds tend to stabilize at values $\sim1-5\,{\rm km\,s^{-1}\,pc^{-1}}$ .
The process is analogous to the well-studied process of bar slowdown in unstable disks (although obviously with the bulge replacing the halo, which is dynamically irrelevant here), and we refer to those studies for further details (????),, although there can be a non- contribution to slowdown from the motion of the BH itself (damping via scattering off the background stars), an effect not included in analytic treatment (comparee.g.??)..
The process is analogous to the well-studied process of bar slowdown in unstable disks (although obviously with the bulge replacing the halo, which is dynamically irrelevant here), and we refer to those studies for further details \citep{weinberg:bar.dynfric, hernquistweinberg92,athanassoula:bar.slowdown, martinezvalpuesta:recurrent.buckling}, although there can be a non-trivial contribution to slowdown from the motion of the BH itself (damping via scattering off the background stars), an effect not included in analytic treatment \citep[compare e.g.][]{adams89:eccentric.instab.in.keplerian.disks, shu:gas.disk.bar.tscale}.
As €) decreases, the barrier in Equation 19 likewise decreases, allowing the mode to strengthen and propagate inwards.
As $\Omega_{p}$ decreases, the barrier in Equation \ref{eqn:slowmode.dispersion.stellar} likewise decreases, allowing the mode to strengthen and propagate inwards.
Second, the mode generates substantial gas inflows, which have three important effects.
Second, the mode generates substantial gas inflows, which have three important effects.
They lead to both gas mass available for further inflow, and form stars, which (discussed below) can sustain a long-lived mode at each radius.
They lead to both gas mass available for further inflow, and form stars, which (discussed below) can sustain a long-lived mode at each radius.
They directly move the fa~0.1 radius inwards, allowing the simple instability region to come closer to the BH.
They directly move the $\tilde{f}_{d}\sim0.1$ radius inwards, allowing the simple instability region to come closer to the BH.
They also, correspondingly, steepen the surface density profile.
They also, correspondingly, steepen the surface density profile.
Inflows will continuously increase the slope as long as inflow is "stalled" at any radius, and this is clearly reflected in Figure 9 as the enclosed disk mass at R«Rerit rises rapidly with time.
Inflows will continuously increase the slope as long as inflow is “stalled” at any radius, and this is clearly reflected in Figure \ref{fig:nuclear.mode.origins} as the enclosed disk mass at $R<R_{\rm crit}$ rises rapidly with time.
At this point, because of significant eccentricity at all radii, v,ο”V. and Κι~ Ro, and orbit crossings can occur at all radii where the mode is supported.
At this point, because of significant eccentricity at all radii, $v_{r}\sim V_{c}$ and $R_{1}\sim R_{0}$ , and orbit crossings can occur at all radii where the mode is supported.
Thus, strong inflows can be sustained in systems such as that in Figure 9 down to radii where the BH completely dominates the potential.
Thus, strong inflows can be sustained in systems such as that in Figure \ref{fig:nuclear.mode.origins} down to radii where the BH completely dominates the potential.
The system quickly reaches aquasi-equilibrium state.
The system quickly reaches aquasi-equilibrium state.
The
The
|Keywords: Pavametric normal form: Hopl-Zero singularity: Eulerian vector fields.
: Parametric normal form; Hopf-Zero singularity; Eulerian vector fields.
(95,6-0.0.
G328.6-0.0.
Comparing the aud coutimmun. we found au sshell around the rregion ROW 91 which indicates that the iregion is enibedded im a molecular cloud.
Comparing the and continuum, we found an shell around the region RCW 94 which indicates that the region is embedded in a molecular cloud.
Iu this case we see the reciprocal effects of massive stars and the siuroundine dduriug the stellar lifetime.
In this case we see the reciprocal effects of massive stars and the surrounding during the stellar lifetime.
The coutimluii ποΓΙ morplology of the rreelontloselv imiatches the iiorphologv of the surroundingr1.
The continuum emission morphology of the region closely matches the morphology of the surrounding.
We use aalbsorpion towards the SNRs C327.0.1 aud 20021.0 to determine kinenatic distances of [23 aud 1.9 kpe. respectively.
We use absorption towards the SNRs G327.4+0.4 and G330.2+1.0 to determine kinematic distances of 4.3 and 4.9 kpc, respectively.
aat the systemic velocity of these τοῖς shows morphological simularitics to f16 contiunuunun enission.
at the systemic velocity of these remnants shows morphological similarities to the continuum emission.
m particular. density cublancemeits were found exterior to reeious of continui liub-brightening for C3:21.1|0.L.
In particular, density enhancements were found exterior to regions of continuum limb-brightening for G327.4+0.4.
We also found two small sshells with no counterparts i1 continuuni euission.
We also found two small shells with no counterparts in continuum emission.
We use the sizes aud lack of detecable expansion velocity to interpret these structures as stalled superuova or wiud blown shells which are older than the radiative lifetimes of either rregious or SNRs.
We use the sizes and lack of detectable expansion velocity to interpret these structures as stalled supernova or wind blown shells which are older than the radiative lifetimes of either regions or SNRs.
Deciphering sstructure las always been crallengine. but the recent availability of Ligh resolution Galactic survevs such as the SGPS has improved the situaion dramatically.
Deciphering structure has always been challenging, but the recent availability of high resolution Galactic surveys such as the SGPS has improved the situation dramatically.
Mich of the inner Galaxy is completely filled with a variety of sstructures includiug shells. worms. sheets. aud filameuts.
Much of the inner Galaxy is completely filled with a variety of structures including shells, worms, sheets, and filaments.
"Though it is extremely dificut to determine the origius of many of the structures wsine cenussion data alone. combinaion with absorption and radio contin emission nieasurenments enables us to deteruue a threc-cimnenusioual. ανασα] picture of the ISM.
Though it is extremely difficult to determine the origins of many of the structures using emission data alone, combination with absorption and radio continuum emission measurements enables us to determine a three-dimensional, dynamical picture of the ISM.
We thank Veta Avedisova for supplving us with her extensive catalogue of star formation regions.
We thank Veta Avedisova for supplying us with her extensive catalogue of star formation regions.
This research las made use of the CDS SI\IBAD database.
This research has made use of the CDS SIMBAD database.
JAID aud NAIALCG acknowledge support of NSF eraut AST-9732695 to the University of Minnesota.
JMD and NMM-G acknowledge support of NSF grant AST-9732695 to the University of Minnesota.
ιο. supported by NAÀSA Cracduate Student Researchers Programs (GSRP) Fellowship NGT 5-502500.
NMM-G is supported by NASA Graduate Student Researchers Program (GSRP) Fellowship NGT 5-50250.
DMCG acknowledecs the support of NASA through Ihble fellowship eraut IIST-ITE-01107.01-À. awarded by STScI. which is operated by AURA Inc. for NASA under couract NAS 5-26555.
BMG acknowledges the support of NASA through Hubble fellowship grant HST-HF-01107.01-A awarded by STScI, which is operated by AURA Inc. for NASA under contract NAS 5-26555.
Ty, we defined the product of the planet/star area ratio and the ratio of the bandpass-integrated planetary to stellar surface photon fluxes, corrected fortransmission?,, to be equal to the measured occultation depth (e.g., Charbonneau et al.
$T_{H}$, we defined the product of the planet/star area ratio and the ratio of the bandpass-integrated planetary to stellar surface photon fluxes, corrected for, to be equal to the measured occultation depth (e.g., Charbonneau et al.
2005).
2005).
We assumed the planet to emit as a black body and, for the star, we used a model spectrum of a G5V star (Pickles 1998), normalised to reproduce the integrated flux of a black body with = 5500KK (Hebb et al.
We assumed the planet to emit as a black body and, for the star, we used a model spectrum of a G5V star (Pickles 1998), normalised to reproduce the integrated flux of a black body with $ = 5500$ K (Hebb et al.
2010).
2010).
The uncertainty in Ty only takes into account the uncertainty in the measured occultation depth.
The uncertainty in $T_{H}$ only takes into account the uncertainty in the measured occultation depth.
1.253 WD|M pairs that have so far been studied. we have provided the solution that the origin of high magnetic fields in white dwarfs relies on a magnetic dvnamo during the common envelope phase of binary. evolution.
$1{,}253$ WD+M pairs that have so far been studied, we have provided the solution that the origin of high magnetic fields in white dwarfs relies on a magnetic dynamo during the common envelope phase of binary evolution.
Svstems with the strongest magnetic. fields emerge from. the CIS. phase either as merged. single stars or with their secondary. stars more nearly in contact with their Roche lobes thus reducing their chance of being detected as pre-magnetic cataclysmic variables.
Systems with the strongest magnetic fields emerge from the CE phase either as merged single stars or with their secondary stars more nearly in contact with their Roche lobes thus reducing their chance of being detected as pre-magnetic cataclysmic variables.
Our hypothesis also predicts the existence of high field magnetic white dwarls that result from systems that merge during the common envelope phase.
Our hypothesis also predicts the existence of high field magnetic white dwarfs that result from systems that merge during the common envelope phase.
“Phe absence. of any AICVs in detached binary stars leacs us to conclude that all highlv magnetic white dwarls have formed in this wav.
The absence of any MCVs in detached binary stars leads us to conclude that all highly magnetic white dwarfs have formed in this way.
We thank Gary Schmidt anc Adam Burrows for enlightching discussions.
We thank Gary Schmidt and Adam Burrows for enlightening discussions.
CAT thanks Churchill College for a Fellowship.
CAT thanks Churchill College for a Fellowship.
DTW is grateful to the Institute of Astronomy. Cambridge for hospitality.
DTW is grateful to the Institute of Astronomy, Cambridge for hospitality.
JL acknowledges support. from. the U.S. National Science Foundation. grant AST 03-07321. for work on white chwarls in the Sloan Digital Sky Survey.
JL acknowledges support from the U.S. National Science Foundation, grant AST 03-07321, for work on white dwarfs in the Sloan Digital Sky Survey.
of the entire SSS population of M31 as viewed by (Di Stefano et al.
of the entire SSS population of M31 as viewed by (Di Stefano et al.
2002b. 2003).
2002b, 2003).
It was clear. however. that not all of the sources identified by these criteria were equally good candidates for the class.
It was clear, however, that not all of the sources identified by these criteria were equally good candidates for the class.
We therefore began by developing strict selection conditions Kong 2003 a).
We therefore began by developing strict selection conditions Kong 2003 a).
These criteria. the so-called HR (hardness ratio) criteria. are the first set of conditions applied in the algorithm presented here.
These criteria, the so-called HR (hardness ratio) criteria, are the first set of conditions applied in the algorithm presented here.
The full algorithm we present here allows us to identify SSSs that are. e.g.. at the high-T end (usually taken to be ~100 eV) of the class. even if they are highly absorbed.
The full algorithm we present here allows us to identify SSSs that are, e.g., at the high-T end (usually taken to be $\sim 100$ eV) of the class, even if they are highly absorbed.
It also covers cases in which a small fraction of the emitted radiation is reprocessed and reemitted in the form of photons of higher energy.
It also covers cases in which a small fraction of the emitted radiation is reprocessed and reemitted in the form of photons of higher energy.
In addition the full algorithm allows us to identify sources that may be somewhat hotter than "elassical" SSSs.
In addition the full algorithm allows us to identify sources that may be somewhat hotter than “classical" SSSs.
This is the first paper to present the full algorithm. and the companion paper tests its efficacy by studying the results of applying it to data from 4 galaxies Kong 2003 b).
This is the first paper to present the full algorithm, and the companion paper tests its efficacy by studying the results of applying it to data from $4$ galaxies Kong 2003 b).
It has also been applied to data from M31 eet al.
It has also been applied to data from M31 et al.
2003 a). and from M104 eet al.
2003 a), and from M104 et al.
2003 b).
2003 b).
In $2 we give an overview of physical models for VSSs. the class of source that motivated this work.
In 2 we give an overview of physical models for VSSs, the class of source that motivated this work.
In 83 we provide a phenomenological definition of SSSs designed to help select sources with the physical characteristics discussed in 82.
In 3 we provide a phenomenological definition of SSSs designed to help select sources with the physical characteristics discussed in 2.
In $4 we outline and test a sequential set of selection criteria that can be applied algorithmically to select SSSs as well as sources that are slightly harder (QSSs). in external galaxies.
In 4 we outline and test a sequential set of selection criteria that can be applied algorithmically to select SSSs as well as sources that are slightly harder (QSSs), in external galaxies.