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tvpeL as described by Martinez-Sxkoraetal.(2009a).. | type, as described by \citet{Martinez-Sykora:2009kl}. |
However. in addition. we also find jet-like structures that resemble spicules of (vpe (Alartinez-Svkora2010).. | However, in addition, we also find jet-like structures that resemble spicules of type \citep{Martinez-Sykora:2010vl}. |
These dilfer markedly in their physical characteristics [rom the pervasive (wpeI. | These differ markedly in their physical characteristics from the pervasive type. |
einulation Bl produced on (he order of 100 spicules of (wpeL | Simulation B1 produced on the order of 100 spicules of type. |
These spicules were found to result. from upwardly propagating shocks passing through the upper chromosphere and into the transition region aud corona. | These spicules were found to result from upwardly propagating shocks passing through the upper chromosphere and into the transition region and corona. |
The shock-driven jets show up as narrow linear structures reaching some few Ali above (the ambient chromosphere. rising and falling with velocities of order 30 km !. and having lifetimes of order 3-5 minutes. | The shock-driven jets show up as narrow linear structures reaching some few Mm above the ambient chromosphere, rising and falling with velocities of order 30 km $^{-1}$, and having lifetimes of order 3-5 minutes. |
An example of thiV. (vpe of jet is seen on the right-hand sile of Fig. 1.. ( | An example of this type of jet is seen on the right-hand side of Fig. \ref{fig:init}, , ( |
blue arrow near x=14 Mm) which appears like a narrow spike in the transition region. | blue arrow near $x=14$ Mm) which appears like a narrow spike in the transition region. |
On the other hand. the simulation produced (wo jets that resemble spicules of (wpe11. but only (wo of them: one of these iV. shown in Fie. | On the other hand, the simulation produced two jets that resemble spicules of type, but only two of them; one of these is shown in Fig. |
1 asa prominent feature in (he transition region (blue arrow near =7 Mm). | \ref{fig:init} as a prominent feature in the transition region (blue arrow near $x=7$ Mm). |
In this latter (vpe of jet. chromospheric material is ejected far into the corona al very. hie[un velocity and heated. reaching velocities up to 95 km I: the jet appears to fade without falling back towrds the solar surface. | In this latter type of jet, chromospheric material is ejected far into the corona at very high velocity and heated, reaching velocities up to 95 km $^{-1}$; the jet appears to fade without falling back towards the solar surface. |
Both of the candidate spicules of type are located near the footpoints of coronal loops. | Both of the candidate spicules of type are located near the footpoints of coronal loops. |
Therefore. the field lines along which the jet is ejected are either open or. at least. penetrate far into the corona. | Therefore, the field lines along which the jet is ejected are either open or, at least, penetrate far into the corona. |
Moreover. both jets appear near «=7 Mm (see blue arrow in Fig 1)). close (o locations where flux emergence is vigorous. | Moreover, both jets appear near $x=7$ Mm (see blue arrow in Fig \ref{fig:init}) ), close to locations where flux emergence is vigorous. |
In fact. we note that both jets first appear alter (he rising [lux tube injected at the bottom boundary. has emerged to and. in part. has crossed the photosphere. | In fact, we note that both jets first appear after the rising flux tube injected at the bottom boundary has emerged to and, in part, has crossed the photosphere. |
In Fig. 2.. | In Fig. \ref{fig:field}, |
details of the (vpe spicule structure are shown for the jet that is active at lime /= 1850s in the D1 simulation: in the figure. isosurfaces for the upflow velocity (blue). temperature (orange: transition region: red: hot corona). and Joule heating per particle (Green) are shown. | details of the type spicule structure are shown for the jet that is active at time $t=1850$ s in the B1 simulation: in the figure, isosurfaces for the upflow velocity (blue), temperature (orange: transition region; red: hot corona), and Joule heating per particle (green) are shown. |
The grevscale map towards the bottom of the figure corresponds (o the horizontal magnetic field strength in the photosphere which clearly shows where flux emergence is vigorous. | The greyscale map towards the bottom of the figure corresponds to the horizontal magnetic field strength in the photosphere which clearly shows where flux emergence is vigorous. |
In addition. the figure shows magnetic field lines close to the jet structure. | In addition, the figure shows magnetic field lines close to the jet structure. |
The jet is formed as chromospheric material isejected into the corona aud heated. with velocities reaching up to 95 km |. | The jet is formed as chromospheric material isejected into the corona and heated, with velocities reaching up to $95$ km $^{-1}$. |
It is composed of many. elements that together comprise a complicated phenomenon. | It is composed of many elements that together comprise a complicated phenomenon. |
In (his paper. the spicule is identified with ejected malerialof chromospheric density. and temperature. the jet-like structure seen in the | In this paper, the spicule is identified with ejected materialof chromospheric density and temperature, the jet-like structure seen in the |
Camuna-ray lines cuited ii solar flares proxride information ou the composition. spectrum. aud aneular distribution accelerated jous flat luteract in the solar atuosphere (Raimaty.ho-zlovaky.&Lingentelter 1979). | Gamma-ray lines emitted in solar flares provide information on the composition, spectrum, and angular distribution of accelerated ions that interact in the solar atmosphere \citep{rkl79}. |
. They also provide information on the compition and deusitv of he ο ΠΕη Ramat | They also provide information on the composition and density of the ambient plasma. |
y&Crauueie976) performed the first calculations of the spectral shape of emission li1ο due to Do»pler shifts froui recoll of excited nuclei after mnupact bv proous and o particles. | \citet{ram76} performed the first calculations of the spectral shape of emission lines due to Doppler shifts from recoil of excited nuclei after impact by protons and $\alpha$ particles. |
They found that the 19O line at 6.13 MeV. has a width of ~100 τον (EWIIM) under bombardiieut by an isotropic distribuion of protous. | They found that the $^{16}$ O line at 6.13 MeV has a width of $\sim 100$ keV (FWHM) under bombardment by an isotropic distribution of protons. |
This width increases to ~160 keV. under bombardment bv a particles. | This width increases to $\sim 160$ keV under bombardment by $\alpha$ particles. |
They also estimated the shift of the line centroid. for a flare rear the center of the solar disk. for both a downward isotropic distriution and a dowmward isotropic distribution limited to angles <30«ee from he radial direction. | They also estimated the shift of the line centroid, for a flare near the center of the solar disk, for both a downward isotropic distribution and a downward isotropic distribution limited to angles $\le 30\deg$ from the radial direction. |
These estimates indicated rec shifts of ~25 aud 0 keV or proton bonibardiueut awl ~{5 and 60 keV for a-particle bonharcdincut. respecively. | These estimates indicated red shifts of $\sim 25 $ and 40 keV for proton bombardment and $\sim 45$ and 60 keV for $\alpha$ -particle bombardment, respectively. |
Rauaty.EKozlovskw.& Liugoaiuelter(1979) xerforiied dezüled calculaions of the shapes aud shifts of eanuna-rav lines YOU P C ane 160 that ive jeco1ue the basis of a \oute €Tarlo code tundamental to several sibseqcut puications. | \citet{rkl79} performed detailed calculations of the shapes and shifts of gamma-ray lines from $^{12}$ C and $^{16}$ O that have become the basis of a Monte Carlo code fundamental to several subsequent publications. |
Aburphns.dKozlovsky.&Ramaty(1988)| used updated cross sections in this code to caculate the solu - “Be. LU*Li. and °C Aline profiles∙ for: isotropic. fan beam. and dowuwiud heaned proton and a-particle distributions having Besse fiction euergv spectra. | \citet{murphy88} used updated cross sections in this code to calculate the solar $^7$ Be, $^7$ Li, and $^{12}$ C line profiles for isotropic, fan beam, and downward beamed proton and $\alpha$ -particle distributions having Bessel function energy spectra. |
They also inchκ). contributions to the PC line from spallation of MO. They described how these angular distijbutiouns cau be | They also included contributions to the $^{12}$ C line from spallation of $^{16}$ O. They described how these angular distributions can be |
This star formation history matches most mcasured UV.continua aud Ito Imuinositv densities. aud includes an upsvard correction for dust reddening of 4390=1.2 mag. | This star formation history matches most measured UV–continuum and $\alpha$ luminosity densities, and includes an upward correction for dust reddening of $A_{1500}=1.2$ mag. |
The SER increases rapidly between :=0 and :=LI. peaks betweeu >=] and 2. aud gently dechues at higher redshifts. | The SFR increases rapidly between $z=0$ and $z=1$, peaks between $z=1$ and 2, and gently declines at higher redshifts. |
Because of the mncertaitics associated with the icoupleteness of the data sets and the amount of dust extinction at carly epochs. we cousider a second scenario m which the SFR remains instead roughly constant at ~2=z (Steidel 1999). (SE2). | Because of the uncertainties associated with the incompleteness of the data sets and the amount of dust extinction at early epochs, we consider a second scenario in which the SFR remains instead roughly constant at $z\gta
2$ (Steidel 1999), (SF2). |
Some recent studies have sugeested that the evolution of the SFR up to 2%1 max have been overestimated (0.8. Cowie. Songaila. Barger 1999). while the rates at + may have been severely underestimated due to large amounts of dust extinction (e.g. Blain 1999). | Some recent studies have suggested that the evolution of the SFR up to $z\approx 1$ may have been overestimated (e.g. Cowie, Songaila, Barger 1999), while the rates at $z$ may have been severely underestimated due to large amounts of dust extinction (e.g. Blain 1999). |
We then consider a third SFR.15M. (SE3). with more star formation at carly epochs. | We then consider a third SFR, (SF3), with more star formation at early epochs. |
In every case we adopt a Salpeter mitial mass fuuctiou (IME) asstuned to remin constant with fime — with a lower cutoff around 0.5AL. (Aladau Pozzetti 2000). consistent with observations of M subdwart disk stars (Could. Bahcall. Εν 1996). | In every case we adopt a Salpeter initial mass function (IMF) – assumed to remain constant with time – with a lower cutoff around $0.5\,M_\odot$ (Madau Pozzetti 2000), consistent with observations of M subdwarf disk stars (Gould, Bahcall, Flynn 1996). |
A coustaut multiplicative factor of 1.67 will convert the SFR to a Salpeter IAIF with a cutoff of 0.1... | A constant multiplicative factor of 1.67 will convert the SFR to a Salpeter IMF with a cutoff of $0.1\,M_\odot$. |
To iuclude the effect of a A dominated cosmology we have computed the difference in luminosity density between an EdS anda A universe. aud applied this correction to the SER above (see Appendix. A for details), | To include the effect of a $\Lambda$ –dominated cosmology we have computed the difference in luminosity density between an EdS and a $\Lambda$ universe, and applied this correction to the SFR above (see Appendix A for details). |
Asstunine that all stars with masses M>8M. explode as corecollapse supernovac (SNe). the SN II rate density Ztax(:) can then be estimated by multiplying the selected SFR by the coefficieut ου... where ο} is the INE and a the stellar mass in solar units. | Assuming that all stars with masses $M>8\,{\rm M}_\odot$ explode as core–collapse supernovae (SNe), the SN II rate density $R_{\rm SN}
(z)$ can then be estimated by multiplying the selected SFR by the coefficient = 0.0122, where $\phi(m)$ is the IMF and $m$ the stellar mass in solar units. |
The resulting rates agree within the errors with the locally observed value of ux=(1.10.1)«10EDSvroDMpe? (e.g. Madan. della Valle. Panagia 1998 ane references therein). | The resulting rates agree within the errors with the locally observed value of $R_{\rm
SN}=(1.1 \pm 0.4) \times 10^{-4}\,h_{65}^3 \,{\rm yr}^{-1}\, {\rm Mpc}^{-3}$ (e.g. Madau, della Valle, Panagia 1998 and references therein). |
The observed fiuxes from GRBs with secure redshifts rule out the classical standard candle hvpothesis (see. Table 1 of Lamb Reichart 2000 and references thereiu): the inferred ‘isotropicequivalent’ photon huuiuosities at peak vary by about a factor of 50. with a mean value of 38h42.«10s1, | The observed fluxes from GRBs with secure redshifts rule out the classical standard candle hypothesis (see Table 1 of Lamb Reichart 2000 and references therein): the inferred `isotropic–equivalent' photon luminosities at peak vary by about a factor of 50, with a mean value of $3.8\, h_{65}^{-2} \times 10^{58} \,{\rm s}^{-1}$. |
The data are too sparse, however. for an empirical determination of the burst luminosity function. CL). | The data are too sparse, however, for an empirical determination of the burst luminosity function, $\psi(L)$. |
To model the ummber counts we then simply assume that the burst huuinosity distribution does uot evolve with redshift and adopt a simple functional form for o£). πα-€ where L denotes the peak Iuninositv iu the 302000 keV cucrev range (restframe). 5is the asviuptotic slope at the bright eud. £o marks a characteristic cutoff scale. aud the coustaut C—[£Z9I(.>1)]| (for ~ <1) ensures a proper normalization f,"oCLyd= 1. | To model the number counts we then simply assume that the burst luminosity distribution does not evolve with redshift and adopt a simple functional form for $\psi(L)$, (L)=C ), where $L$ denotes the peak luminosity in the 30–2000 keV energy range (rest–frame), $\gamma$is the asymptotic slope at the bright end, $L_0$ marks a characteristic cutoff scale, and the constant $C=[L_0\Gamma(-\gamma-1)]^{-1}$ (for $\gamma<-1$ ) ensures a proper normalization $\int_0^\infty \psi(L) dL=1$ . |
concentration of ος=6.4. which is about 16 per cent vigher than the concentrations found. for normal groups fy,=5.5. | concentration of $c_{1/5} = 6.4$, which is about 16 per cent higher than the concentrations found for normal groups $c_{1/5} =
5.5$. |
Our findings are consistent with fossil. groups ong systems with higher dark matter concentrations than usual groups. which supports such a trend. found. hy ?.. | Our findings are consistent with fossil groups being systems with higher dark matter concentrations than usual groups, which supports such a trend found by \citet[][]{Khosroshahi2007a}. |
At present there are νο few observational constraints on his issue (e.g.77).. | At present there are yet few observational constraints on this issue \cite[e.g.][]{Gastaldello2007a,
Khosroshahi2007a}. |
Llowever with the upcoming X-ray observations of fossil groups with Chandra and NMM will oovide soon better constraints. | However with the upcoming X-ray observations of fossil groups with Chandra and XMM will provide soon better constraints. |
ltecent studies of the giant elliptical at the center of fossil eroups report no signs of a recent major merger activity. indicating that any major merger should have happened at least more than approximately 3 Civrs ago (?2).. | Recent studies of the giant elliptical at the center of fossil groups report no signs of a recent major merger activity, indicating that any major merger should have happened at least more than approximately 3 Gyrs ago \cite[][]{Jones2000a, Khosroshahi2006a}. |
For cach halo we estimate the time of the last. major merger of the group haloes of our sample by studying the detailed mass assembly. history. | For each halo we estimate the time of the last major merger of the group haloes of our sample by studying the detailed mass assembly history. |
To identify the time of the last major merger. we denote a halo as a major merger remnant if its major progenitors were classified: as a single group at one time but two separate groups with a mass ratio less than 41 at the preceeding time. | To identify the time of the last major merger, we denote a halo as a major merger remnant if its major progenitors were classified as a single group at one time but two separate groups with a mass ratio less than 4:1 at the preceeding time. |
Note that when the mass ratio defining the major merger event is restricted to almost 1:1 progenitors. the time of the last major merger should in general coincide with the formation time (as defined above). | Note that when the mass ratio defining the major merger event is restricted to almost 1:1 progenitors, the time of the last major merger should in general coincide with the formation time (as defined above). |
We find that only of the fossil groups experienced the last major merger less than 2 Civrs ago. and at least had the last major merger longer than 6 Cives in qualitative agreement with the observations. | We find that only of the fossil groups experienced the last major merger less than 2 Gyrs ago, and at least had the last major merger longer than 6 Gyrs in qualitative agreement with the observations. |
ln previous sections we investigated. some properties of our galaxv-group sized haloes that can be. compared. to | In previous sections we investigated some properties of our galaxy-group sized haloes that can be compared to |
[orbidden emission line radiation. | forbidden emission line radiation. |
Ássuming a jet velocity of 200 km/s (which is the average jel velocity. lor CTTSs derived by Hirthetal.(1994))) ancl adopting a radial velocity of - 20 km/s for the [OI] A6300 peaks. we derive an inclination angle for the jet to the line of sight of 84°. | Assuming a jet velocity of 200 km/s (which is the average jet velocity for CTTSs derived by \citet{Hirth94}) ) and adopting a radial velocity of $\pm$ 20 km/s for the [OI] $\lambda$ 6300 peaks, we derive an inclination angle for the jet to the line of sight of $\degr$. |
This is consistent with a general picture for the svstem in which we view the XT star close to the orbital plane of its circumstellar (or circumbinary) disk aud the jets emerge roughly perpendicular (ο the disk plane. as for the star associated with HII 30 and other examples imaged by the IIubble Space Telescope (Ravetal.1996). | This is consistent with a general picture for the system in which we view the K7 star close to the orbital plane of its circumstellar (or circumbinary) disk and the jets emerge roughly perpendicular to the disk plane, as for the star associated with HH 30 and other examples imaged by the Hubble Space Telescope \citep{Ray96}. |
. The absence of anv significant variation in the profiles or flux of the forbidden line radiation during eclipse is consistent with the expectation. based on the CTTS analogy. that it arises at distances of tens of AU's from the star. bevond the region variably occultec. | The absence of any significant variation in the profiles or flux of the forbidden line radiation during eclipse is consistent with the expectation, based on the CTTS analogy, that it arises at distances of tens of AU's from the star, beyond the region variably occulted. |
The behavior of the hydrogen lines is complex because some components do arise close to the star and. therefore. suffer variable occultation effects along with the photosphere. | The behavior of the hydrogen lines is complex because some components do arise close to the star and, therefore, suffer variable occultation effects along with the photosphere. |
Since we expect that the Ila line has a much higher optical depth (han the IL? line. it is obvious that anv emission line region will be more extended in Ila. than Ilo. | Since we expect that the $\alpha$ line has a much higher optical depth than the $\beta$ line, it is obvious that any emission line region will be more extended in $\alpha$ than $\beta$. |
If. we assume that the nagnelic axis of (he INT star is (ilted toward us at 5-107. which is a reasonable assumption supported by the [OI] line profiles. at mid-eclipse. one can qualitatively understand the Ha ine profile as resulting from low velocity material in (he outer. more extended La emission region while the star and the Ilo line forming region are obscured by the oceulting disk naterial. | If we assume that the magnetic axis of the K7 star is tilted toward us at $\sim$ $\degr$, which is a reasonable assumption supported by the [OI] line profiles, at mid-eclipse, one can qualitatively understand the $\alpha$ line profile as resulting from low velocity material in the outer, more extended $\alpha$ emission region while the star and the $\beta$ line forming region are obscured by the occulting disk material. |
This would also explain why there is almost no IL? flux curing mid-eclipse. | This would also explain why there is almost no $\beta$ flux during mid-eclipse. |
During eeress. we expect most of the [la emission line region. as well as some of the IL? emission ine reelon. which is closer to (he star. to be visible. | During egress, we expect most of the $\alpha$ emission line region, as well as some of the $\beta$ emission line region, which is closer to the star, to be visible. |
The IIa line profile during egress has a ow velocitv emission peak with a central absorption feature similar to what is seen in (he out-of-eclipse profile. | The $\alpha$ line profile during egress has a low velocity emission peak with a central absorption feature similar to what is seen in the out-of-eclipse profile. |
AddiGionally. two "shoulders" appear along the profile at about x 150 kin/s extending out to z 300 km/s. These "shoulders" could be due to material rotating; in the outer parts of (he magnetosphere. | Additionally, two “shoulders” appear along the profile at about $\pm$ 150 km/s extending out to $\pm$ 300 km/s. These “shoulders” could be due to material rotating in the outer parts of the magnetosphere. |
However. given that the e1Η(1) is measured to be <5 km/s. and that the magnetosphere only extends out to about 5-6 stellar radii. (le maximum rotational velocitv is about 5-6 esin() or 25-30 km/s. makine it diflieult to explain the La enussion-line profile with a rotating magnetopshere. | However, given that the $v\thinspace sin(i)$ is measured to be $<$ 5 km/s, and that the magnetosphere only extends out to about 5-6 stellar radii, the maximum rotational velocity is about 5-6 $v\thinspace sin(i)$ or 25-30 km/s, making it difficult to explain the $\alpha$ emission-line profile with a rotating magnetopshere. |
A more attractive hypothesis is that the high velocity "shoulders" on the Πα line. so prominent in the egress spectrum. arise [rom material falling along magnetic accretion columus. | A more attractive hypothesis is that the high velocity “shoulders” on the $\alpha$ line, so prominent in the egress spectrum, arise from material falling along magnetic accretion columns. |
This interpretation can also qualitatively account for the IL? emission line prolile during egress. | This interpretation can also qualitatively account for the $\beta$ emission line profile during egress. |
Aclopling values for the mass and radius of the T star from Table 1 of l]Luniltonetal.(2001). a [ree-Lall velocity of about 380 km/s can be associated with malerial al the surface of the star. | Adopting values for the mass and radius of the K7 star from Table 1 of \citet{Ham01}, a free-fall velocity of about 380 km/s can be associated with material at the surface of the star. |
Since IL2 is produced much closer to Cie star. the wing extending to nearly -350 km/s could be representative of material accreting along magnetic field lines near the pole. | Since $\beta$ is produced much closer to the star, the blue-shifted wing extending to nearly -350 km/s could be representative of material accreting along magnetic field lines near the pole. |
The asymmetry seen in the IL2 emission line prolile is most likely due to the [act that the star is slightly inclined toward our line of sight. | The asymmetry seen in the $\beta$ emission line profile is most likely due to the fact that the star is slightly inclined toward our line of sight. |
The | The |
The sol X-ray (ransicnts (ονTs) are a subclass of the lowmass Xrav binaries in which a lowmass star transfers luaerial to a neutron star or black hole. | The soft X-ray transients (SXTs) are a subclass of the low--mass X–ray binaries in which a low–mass star transfers material to a neutron star or black hole. |
“Phev undergo large otbursts. reaching X.ταν luminosities of order the Iclineton limit (sce Tanaka Lewin 1995 and Tanaka Shixvzaki 1996 for reviews). | They undergo large outbursts, reaching X–ray luminosities of order the Eddington limit (see Tanaka Lewin 1995 and Tanaka Shibazaki 1996 for reviews). |
Phe outbursts have similarities tothose of dwarf novae. but with important dilferences. one being that the timescales are much longer: a dwarf nova oulrust lasts a few days and typically recurs every few weeks. whereas in ονΕς the corresponding timescales are months and vears. | The outbursts have similarities to those of dwarf novae, but with important differences, one being that the timescales are much longer; a dwarf nova outburst lasts a few days and typically recurs every few weeks, whereas in SXTs the corresponding timescales are months and years. |
Lt is generally accepted that the dwarf nova outburst results from a thermalviscous instability in the aceretion disc (see. Cannizzo 1993 for a review). | It is generally accepted that the dwarf nova outburst results from a thermal–viscous instability in the accretion disc (see Cannizzo 1993 for a review). |
‘Lo some extent the models used for dwarf novae can reproduce the correct. timescales for the ο outburst. but only by reducing the viscosity. parameter by a factor of 10.100 and choosing a particular functional form for it (e.g. Cannizzo. Chen Livio. 1995). | To some extent the models used for dwarf novae can reproduce the correct timescales for the SXT outburst, but only by reducing the viscosity parameter by a factor of 10–100 and choosing a particular functional form for it (e.g. Cannizzo, Chen Livio, 1995). |
Although there are similarities between the disces during outburst in chwarf novae and SNTs. an important observed cillerencee is that the disces in the SXTs are heavily irraciatec | Although there are similarities between the discs during outburst in dwarf novae and SXTs, an important observed difference is that the discs in the SXTs are heavily irradiated |
Although there are similarities between the disces during outburst in chwarf novae and SNTs. an important observed cillerencee is that the disces in the SXTs are heavily irraciatecἱ | Although there are similarities between the discs during outburst in dwarf novae and SXTs, an important observed difference is that the discs in the SXTs are heavily irradiated |
X-ray. UV. and thermal radio emission (van der Ixreuit de Bruyn 1976. Code Welch 1982. Fabbiano. Feigelson Zamorani 1982. Zezas. Georgantopoulos Ward. 1998. Conselice et al. | X-ray, UV, and thermal radio emission (van der Kruit de Bruyn 1976, Code Welch 1982, Fabbiano, Feigelson Zamorani 1982, Zezas, Georgantopoulos Ward 1998, Conselice et al. |
2000). intense UV. and optical emission lines tvpical of OB associations (Heckman Balick 1980. Kinney et al. | 2000), intense UV and optical emission lines typical of OB associations (Heckman Balick 1980, Kinney et al. |
1993. Mulder ct al. | 1993, Mulder et al. |
1995 and references therein) and its large global Lla equivalent. width (MVDDB95. Mulder van Driel 1996 hereafter MyD96]. and references therein) | 1995 and references therein) and its large global $\alpha$ equivalent width (MvDB95, Mulder van Driel 1996 [hereafter MvD96], and references therein). |
A number of morphological peculiaritics in its outer nuts (eg. Balick Lleekman 1981. Kregel Sancisi 2001. hereafter. BLISL. ΟΙ. ALVD96). combined with the disturbed: kinematies of theH1 gas in the inner regions (c... van der Ixruit. 1976. MvDD95. INSOL). the mismatch oetween stellar and gas-dvnamical geometry. (e$... 1151). and the large number of earlv-tvpe stars required to explain he galaxv's La emission. (van der νο 1976. DIISI). suggest that NGC 3310 was allected by a major gravitational disturbance. | A number of morphological peculiarities in its outer parts (e.g., Balick Heckman 1981, Kregel Sancisi 2001, [hereafter BH81, KS01], MvD96), combined with the disturbed kinematics of the gas in the inner regions (e.g., van der Kruit 1976, MvDB95, KS01), the mismatch between stellar and gas-dynamical geometry (e.g., BH81), and the large number of early-type stars required to explain the galaxy's $\alpha$ emission (van der Kruit 1976, BH81), suggest that NGC 3310 was affected by a major gravitational disturbance. |
This led to high. possibly sustained. star ormation rates in the past LOO Myr (cf | This led to high, possibly sustained, star formation rates in the past $\sim 100$ Myr (cf. |
DIISI). | BH81). |
Since attempts to identify a nearby companion galaxy as the cause for the disruption and the expected subsequent major starburst were unsuccessful (ef | Since attempts to identify a nearby companion galaxy as the cause for the disruption and the expected subsequent major starburst were unsuccessful (cf. |
van cer νι 1976. van der Ixruüit de Bruyn 1976). it was suggested that NGC 3310 accreted a eas-rich. but metal-poor companion galaxy. which subsequently fragmented as a result of the encounter (BIISL. Schweizer Seitzer LOSS. MvDDB95. MVvD96. Smith et al. | van der Kruit 1976, van der Kruit de Bruyn 1976), it was suggested that NGC 3310 accreted a gas-rich, but metal-poor companion galaxy, which subsequently fragmented as a result of the encounter (BH81, Schweizer Seitzer 1988, MvDB95, MvD96, Smith et al. |
1996. hereafter SOG: modeled by Athanassoula 1992 and Piner. Stone “Teuben 1995). or perhaps we are currently seeing a newlv-formed. disc. | 1996, hereafter S96; modeled by Athanassoula 1992 and Piner, Stone Teuben 1995), or perhaps we are currently seeing a newly-formed disc. |
This argument was based predominantly on the absence of any close companion ealaxy and on the unusually low (subsolar) metallicity found in star forming regions surrounding the nucleus. although the nucleus itself. appears to have solar metallicity. (e.g. lleckman Balick 1980. Puxlev. Hawarden Mountain 1990. Pastoriza et al. | This argument was based predominantly on the absence of any close companion galaxy and on the unusually low (subsolar) metallicity found in star forming regions surrounding the nucleus, although the nucleus itself appears to have solar metallicity (e.g., Heckman Balick 1980, Puxley, Hawarden Mountain 1990, Pastoriza et al. |
1993. hereafter P93). | 1993, hereafter P93). |
Additional support for this interpretation is provided by the [fa-ultraviolet and. Lehane 13 ‘luminosity profiles. tvpical of late-stage galaxy mergers (806. WSOL). and the optical and tails and ripples observed in the outer parts. | Additional support for this interpretation is provided by the far-ultraviolet and -band $R^{1/4}$ luminosity profiles, typical of late-stage galaxy mergers (S96, KS01), and the optical and tails and ripples observed in the outer parts. |
The galaxys most conspicuous visual peculiarity is the wellstucdiecl bright optical "bow and arrow” structure (nomenclature first. used. by Walker Chincarini 1967). | The galaxy's most conspicuous visual peculiarity is the well-studied bright optical “bow and arrow” structure (nomenclature first used by Walker Chincarini 1967). |
The carow is a jet-like structure with possibly an counterpart. (although the latter is significantly more extended: AlyDBS05). which has been interpreted. as the result of a nuclear explosive event some Ls.10' vr ago (Bertola Sharp 1984). or. in combination with the ripple pattern that includes the "bow" — as the result of a recent merger with a smaller companion galaxy (or even a third. smaller object: INSOI). | The “arrow” is a jet-like structure with possibly an counterpart (although the latter is significantly more extended; MvDB95), which has been interpreted as the result of a nuclear explosive event some $1.8 \times 10^7$ yr ago (Bertola Sharp 1984), or – in combination with the ripple pattern that includes the “bow” – as the result of a recent merger with a smaller companion galaxy (or even a third, smaller object; KS01). |
The latter explanation seems more likely (see MvDDB95. IXSOD). | The latter explanation seems more likely (see MvDB95, KS01). |
These authors argued. that a “projectile” object hitting the dise of NGC 3310 at a fairly oblique angle could have been drawn out into the observed configuration. which is supported by the anomalous velocity structure observed. within these features (AIvDB95): the optical “arrow” is then best understood. as ai series. of regions in which active star formation was triggered due to the collision. | These authors argued that a “projectile” object hitting the disc of NGC 3310 at a fairly oblique angle could have been drawn out into the observed configuration, which is supported by the anomalous velocity structure observed within these features (MvDB95); the optical “arrow” is then best understood as a series of regions in which active star formation was triggered due to the collision. |
Lf this scenario is correct. the merger event must have occurred. z10 Myr ago to prevent. significant reclistribution of the longer jet due to dillerential rotation (cf. | If this scenario is correct, the merger event must have occurred $\lesssim 10$ Myr ago to prevent significant redistribution of the longer jet due to differential rotation (cf. |
MvDB95). although the outer jet regions appear to curve away [rom the line of sight (INSOL. see also MvDD95). | MvDB95), although the outer jet regions appear to curve away from the line of sight (KS01, see also MvDB95). |
The bar-driven star formation scenario suggested. by Conselice ct al. ( | The bar-driven star formation scenario suggested by Conselice et al. ( |
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