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From this figure it is evident that in some stars HHD 183143 and HD 20041) A6270 is much stronger than A4964. while in other stars HHD 147888 and HD 147889) the situation is reversed. | From this figure it is evident that in some stars HD 183143 and HD 20041) $\lambda$ 6270 is much stronger than $\lambda$ 4964, while in other stars HD 147888 and HD 147889) the situation is reversed. |
This clearly rules out the possibility that both bands can be due to the same carrier. and therefore they cannot both be due toC;. | This clearly rules out the possibility that both bands can be due to the same carrier, and therefore they cannot both be due to. |
. There are two other weak vibronic bands of the A«X transition of tthat were reported by Tule]etal.(1998). | There are two other weak vibronic bands of the $A \leftarrow X$ transition of that were reported by \citet{tulej}. |
. These both happen to be doublets: 17 at 5089.5 and 5095.7A.. and 113], at 5449.6 and 5456.7A. | These both happen to be doublets: $1^2_0$ at 5089.5 and 5095.7, and $1^1_03^1_0$ at 5449.6 and 5456.7. |
. We were not able to detect these bands in our astronomical spectra. but because of the intrinsic weakness of these bands (compared with the origin band) we were not able to set useful upper limits on them either. | We were not able to detect these bands in our astronomical spectra, but because of the intrinsic weakness of these bands (compared with the origin band) we were not able to set useful upper limits on them either. |
Similarly. we were not able to obtain a useful limit for the origin band of the B«Χ | Similarly, we were not able to obtain a useful limit for the origin band of the $B \leftarrow X$ |
Alereer of galaxies is one of the major processes of ealaxy formation in the ierarchical ACDM. cosinogonx. | Merger of galaxies is one of the major processes of galaxy formation in the hierarchical$\Lambda$ CDM cosmogony. |
Galaxy niergers can naπαν lead to the formation of paired and binary nassive black holes (MDITs). since inost galaxies. es]ecallv those with spheroidal components. host MDITs at their ceuters2002). | Galaxy mergers can naturally lead to the formation of paired and binary massive black holes (MBHs), since most galaxies, especially those with spheroidal components, host MBHs at their centers. |
. If both of the mereiue galaxies are gas rich. a larec amouut of eas can be channeled to he ceutral region of eacji nereime galaxw. as indicated No nuuercal simulations1989). | If both of the merging galaxies are gas rich, a large amount of gas can be channeled to the central region of each merging galaxy, as indicated by numerical simulations. |
.. A dual AGN (Δέν) could thei emieree if the accretion onto voth MDIIs is trigecree during the mereie process. | A dual AGN (dAGN) could then emerge if the accretion onto both MBHs is triggered during the merging process. |
Thus the existence of dACNs in the universe aud their deioeraphy provide au inuportaut probe not only of he hierarchical galaxy formation models but also of he trigecringOO mechamisius of nuclear activities aud the assembly history of MDITIs. | Thus the existence of dAGNs in the universe and their demography provide an important probe not only of the hierarchical galaxy formation models but also of the triggering mechanisms of nuclear activities and the assembly history of MBHs. |
Iu the past decade. substautial progress has been uade dn searching for dAGNs and binary MDITS (BBUs) through various possible signatures. such as double-peaked broad liues2009).. double-peaked.. τιarrow lines and various other methods2010). | In the past decade, substantial progress has been made in searching for dAGNs and binary MBHs (BBHs) through various possible signatures, such as double-peaked broad lines, double-peaked narrow lines and various other methods. |
. Following a svetematic scrutiny of the NIR nuages and optical slit spectra of a sample of double-peaked narrow line type 2 ACNs. it has recently beeu reported that roughly of the z<0.3 type 2 AGNs are kpc-cale dACGNs after taking iuto account the selection coupletcucss | Following a systematic scrutiny of the NIR images and optical slit spectra of a sample of double-peaked narrow line type 2 AGNs, it has recently been reported that roughly of the $z\la 0.3$ type 2 AGNs are kpc-scale dAGNs after taking into account the selection completeness. |
This 2011AnaiuDodyCitationEud221|Liu10a.LiuL0b.Sheultü.frequency is surprisingly low even if we were to asstune that all these observations are indeed of dAGNs and not as has been otherwise suggested of superposed single ACNs. bipolar jets and accretion disks. | This frequency is surprisingly low even if we were to assume that all these observations are indeed of dAGNs and not as has been otherwise suggested of superposed single AGNs, bipolar jets and accretion disks. |
The observed frequency of dACIN ids indeed more than an order of magnitude lower than that expected roni the observed iuajor ierecr rate of ealaxies if one makes the assumption that cach major merger would vield a dAGNmaimBodyCitationEnud296|Rosarioll. | The observed frequency of dAGN is indeed more than an order of magnitude lower than that expected from the observed major merger rate of galaxies if one makes the assumption that each major merger would yield a dAGN. |
The low frequency of dACNs. uulikely due to selection effect2011).. seems ecuerally at odds with the commonly accepted scenario that ACGN/QSO. activities are trieecred by major mergers of galaxies. | The low frequency of dAGNs, unlikely due to selection effect, seems generally at odds with the commonly accepted scenario that AGN/QSO activities are triggered by major mergers of galaxies. |
Iu this paper. we construct a phenomenological uodel to address the observed. frequency of kpe-scale dACNs. by taking iuto account the following factors: (1) the ΠΙΟΓΟΟΥ 1ate of galaxies: (2) the types of galaxy mcrecrs that cau trigecr the nuclear activities of both progenitors: (3) when and where significant nuclear activities. are trigeeredoo (bv vsieuificant™ nuclear activities we nean that the Eddington ratios of the nuclear Inninosities are close to l (eg. 011) rather than 10? or less): aud (1) how long the nuclear activities of both progenitors can last before the merger of the two MDIISs. | In this paper, we construct a phenomenological model to address the observed frequency of kpc-scale dAGNs, by taking into account the following factors: (1) the merger rate of galaxies; (2) the types of galaxy mergers that can trigger the nuclear activities of both progenitors; (3) when and where significant nuclear activities are triggered (by “significant” nuclear activities we mean that the Eddington ratios of the nuclear luminosities are close to 1 (e.g., 0.1–1), rather than $10^{-3}$ or less); and (4) how long the nuclear activities of both progenitors can last before the merger of the two MBHs. |
After the morecr of two MDIIS. the uucleus of the merged galaxy i likely to still be active. but it only appears as a sinele AGN. | After the merger of two MBHs, the nucleus of the merged galaxy is likely to still be active, but it only appears as a single AGN. |
In Section ??.. we show how these factors | In Section \ref{sec:model}, , we show how these factors |
Larly-type stars with. high. peculiar. velocities2. (Le.. rüuswvaW stars. with. velocitiesel ος230. kins 1. e.g. mCües Boltou; ↽⋠∖⋅1986) are unconunon. | Early-type stars with high peculiar velocities (i.e. runaway stars, with velocities $v_* > 30$ km $^{-1}$, e.g. Gies Bolton 1986) are uncommon. |
For instance.. Madzz-Apelltuniznan et al. ( | For instance, z-Apellánniz et al. ( |
2001. Galactic O star catalog for V<<8 stars) lists runaway stars out of 370. | 2004, Galactic O star catalog for $V < 8$ stars) lists $\sim$ runaway stars out of 370. |
These particular stars can be identified.n by the perturbation. they produce ii. ambient medium. (e.g.- Nobulnicky et."E al. | These particular stars can be identified by the perturbation they produce in the ambient medium (e.g. Kobulnicky et al. |
2010 timescalesaud refercnees therein). | 2010 and references therein). |
When the strong winds of runaway OD stars sweep relatively large amounts of gas aud dust. he material piles up iu the so-called stellar bow shock. | When the strong winds of runaway OB stars sweep relatively large amounts of gas and dust, the material piles up in the so-called stellar bow shock. |
Bow shocks develop as arc-shaped structures. with bows yolutingo iu the same direction as the stellar velocity. while the star moves supersonicallv iu the surrounding iuterstellar medi (0SMD. | Bow shocks develop as arc-shaped structures, with bows pointing in the same direction as the stellar velocity, while the star moves supersonically in the surrounding interstellar medium (ISM). |
The winds are coufined bv he rin pressure of the ISAL at distances from the star determined by momentum balance. | The winds are confined by the ram pressure of the ISM, at distances from the star determined by momentum balance. |
The stellar aud shock-excited vacation heats the dust aud gas swept by the bow shock. | The stellar and shock-excited radiation heats the dust and gas swept by the bow shock. |
The dust. iu turn. re-raciates the energv as to-far IR excess flux. | The dust, in turn, re-radiates the energy as mid-to-far IR excess flux. |
As soon as IRAS images became available. Van Dureu AlcCray (1988) looked for bow-shaped features near high-velocity O stars (sec also Van Buren et al. | As soon as IRAS images became available, Van Buren McCray (1988) looked for bow-shaped features near high-velocity O stars (see also Van Buren et al. |
1995 ando Noricea-Crespo-: et al. | 1995 and Noriega-Crespo et al. |
41997). | 1997). |
= Thed authors detected an IR candidate close to the Ο supereiaut 133651 ΕΠ — 20x3364077.|IY59'07.107; Lb=κοLL. 12.33%) | The authors detected an IR candidate close to the O supergiant $^{\circ}3654$ $\alpha,\delta$ [J2000] = $20^{\rm h}33^{\rm m}36.077^{\rm s}, +43^{\circ}
59' 07.40''$; $l, b = 82.41^\circ, +2.33^\circ$ ). |
Recently. Comeroun, PasqualiD (2007)50h related the star 1373651 to a bow shock detected with the Midcourse Space. eNperiment~ (MSX)m at D aud E bauds. | Recently, Comerónn Pasquali (2007) related the star $^{\circ}3654$ to a bow shock detected with the Midcourse Space eXperiment (MSX) at D and E bands. |
They studied: the stellar motion relative to the surrounding imaterial . ↴⋅∙↴ ⋅⊲↽⋡↸∖⋯↕↴∖∷∖↴↕∪∐↕↥⋅∪⋯↑∐↸∖↴⋝∪↖↖↽↴∖↴∐⋯⊳↨↘↽∪↕↕≽↕≻↖↓⋅≩⋅≩∢≽⋅↱⊐↓∙ ⋜⋃≺↧↻⊓∏⋯↴∖↸∖≼↧↑∐∖↴∖↑⋜∐↕↴∖⋜↧↥∏∐⋜∏↖⋜↧⋅↖⋯↸∖∐∐⋝↸∖↥↕↥∪⋯≼⋅↖∶↴∙ OD2 association. | They studied the stellar motion relative to the surrounding material and proposed the star is a runaway member from Cyg OB2 association. |
Comerdun auc Pasquali determined a spectral type O LTE and derived an age of about 1.6 Myr and a stell mass of ~ 70 which makes the star one ofd the three more massive: runway stars kuow1 so far. | Comerónn and Pasquali determined a spectral type O4 If, and derived an age of about 1.6 Myr and a stellar mass of $\sim$ 70 $_\odot$, which makes the star one of the three more massive runaway stars known so far. |
On the basis of these estimates. the authors favor a dynamical ejection scenario (see Hooserwerf et al. | On the basis of these estimates, the authors favor a dynamical ejection scenario (see Hoogerwerf et al. |
2000. an2001 forks reviewsget about the originuo of⊳∙ runaway stars). | 2000, 2001 for reviews about the origin of runaway stars). |
Cravanazde. Bomaus; (2008)⊀∖ sugecst mstead⋅ that 13°365 Lis part of a stellar svsteui. formed by a close encounter between two tight massive binaries in the core of Cre OD2. | Gravamazde Bomans (2008) suggest instead that $^{\circ}\,3654$ is part of a stellar system, formed by a close encounter between two tight massive binaries in the core of Cyg OB2. |
The star should be a blue strageler. to match the involvedον in⋅ their⋅ hypothesis. | The star should be a blue straggler to match the timescales involved in their hypothesis. |
⋅ Kobulnicky et. al. he ( | Kobulnicky et al. ( |
2010) measured a heliocentric radial velocity of 00,|9.L ns 1>. and derived. a stellar niass-loss rate ofDM PAD | | 2010) measured a heliocentric radial velocity of $-66.2\pm9.4$ km $^{-1}$, and derived a stellar mass-loss rate of $1.6 \times 10^{-4}$ $_\odot$ $^{-1}$. |
We analyzed data from the NRAO-VLA Sky Survey (NVSS. Condon et al. | We analyzed data from the NRAO-VLA Sky Survey (NVSS, Condon et al. |
1998). | 1998). |
The images revealed a coma-shaped source of ~ 7 arcu. spatially coincident with the AISX structure (see Fieure 1: NVSS aueular resolution: 157:2 uis noise:. ) Landy 1 )). | The images revealed a coma-shaped source of $\sim$ 7 arcmin, spatially coincident with the MSX structure (see Figure 1; NVSS angular resolution: 45”; rms noise: 1 mJy $^{-1}$ ). |
No: point: sources above 5o (l0 1iJv) that are positionally coincideut with the MSX source are detected in the MET-Creeu Bauk Survey (6D6. αι et al. | No point sources above $\sigma$ (40 mJy) that are positionally coincident with the MSX source are detected in the MIT-Green Bank Survey (GB6, Griffith et al. |
1991). | 1991). |
Iuspectiou of the continuum cussion at [08 and 1120 MIIZ with the COPS Survey. (Tavlor et al. | Inspection of the continuum emission at 408 and 1420 MHz with the CGPS Survey (Taylor et al. |
2003. angular resolutions of 3.1 and 1) confirms that he reeion ceutered on Cre OB? is complex aud has strong enission on various augular scales (Peri et al. | 2003, angular resolutions of $3.4'$ and $1'$ ) confirms that the region centered on Cyg OB2 is complex and has strong emission on various angular scales (Peri et al. |
2010). | 2010). |
hA radi UMCLIO8fudGVx οof ththeC 1POW DOWSHOCAhodSIIOCÓS C SICSACU LlielATSoat Q nt1e owesieal rac]processes that give rise to ΟΛΗhieh-energy cussion |Bonn a stellar source. regardless of: the history. of: the runaway star. | A radio study of the bow shock can shed light on the physical processes that give rise to high-energy emission from a stellar source, regardless of the history of the runaway star. |
The shock can accelerate particles up o relativistic‘ enereicso bv Fernü mechauisin. | The shock can accelerate particles up to relativistic energies by Fermi mechanism. |
‘ Encreetico electrons will cool through svuchrotron raciation. xoduciuglnci Dta nonthermalη DUETradio source. | Energetic electrons will cool through synchrotron radiation, producing a nonthermal radio source. |
WiWe carried‘ied outont radio‘. observations at two frequencies to study the nature of the ∙∙≯ BN NEN Tu | We carried out radio observations at two frequencies to study the nature of the emission from the bow shock of $^{\circ}\,3654$. |
this we present the results of the radio 6jbservatious iu the formu: of a spectral iudex map of the Ivow-shock region of the runaway star and distinguish betweenMoo thermalDENM aud. nouthermalWye ciuission∖ regions. | In this we present the results of the radio observations in the form of a spectral index map of the bow-shock region of the runaway star and distinguish between thermal and nonthermal emission regions. |
⋅∖↕↜⋅. W¢7A Ixieflv discuss the issue of whether a bow shock could produce high-cncrev cussion enough to be detected with iustriueuts like Fermi or the future Cherenkov Telescope | We briefly discuss the issue of whether a bow shock could produce high-energy emission enough to be detected with instruments like Fermi or the future Cherenkov Telescope |
Paleoily. distributed over the age of t16 Galaxy and cach procucing an equal umber of brown cawarts over the same formation timescale. | randomly distributed over the age of the Galaxy and each producing an equal number of brown dwarfs over the same formation timescale. |
The birth rate distribution of cach cluster was assumed to be Caussian with a characteristic time scale z4=10 Myr. | The birth rate distribution of each cluster was assumed to be Gaussian with a characteristic time scale ${\tau}_{cl} = 10$ Myr. |
These assmuptious are not necessarily representative of the true vields aud lifetimes of voung clusters in the Galaxy. but are suitable for this study. | These assumptions are not necessarily representative of the true yields and lifetimes of young clusters in the Galaxy, but are suitable for this study. |
Finally. the fifth (vhalo”) birth rate cousiders oulv brown dwarfs bora within a 1 Cr burst 9 Covr in the ost. and is meant to represent f1ο conditions of the Calactic halo or old globular cluster substellar populations (Reid&Tlawley2000). | Finally, the fifth (“halo”) birth rate considers only brown dwarfs born within a 1 Gyr burst 9 Gyr in the past, and is meant to represent the conditions of the Galactic halo or old globular cluster substellar populations \citep{rei00b}. |
. For each of these birth rates. an age range of 0.01.<Ft10 Cir is nominally adopted. although minium ages of 1 to 100 May were also examined to investigate the contribution of voung populations in the simulated LFs (see 6 1.2.3). | For each of these birth rates, an age range of $0.01 \leq t \leq 10$ Gyr is nominally adopted, although minimum ages of 1 to 100 Myr were also examined to investigate the contribution of young populations in the simulated LFs (see $\S$ 4.2.3). |
The choice of [a dnctallicity distribution is primarily constrained by the evolutionary models used (5 2.3). both of which assimnue solar abundances. | The choice of a metallicity distribution is primarily constrained by the evolutionary models used $\S$ 2.3), both of which assume solar abundances. |
Therefore. a constant distribution P(Z) = 1 is adopted with Z=Z... | Therefore, a constant distribution $P(Z)$ = 1 is adopted with $Z = Z_{\sun}$. |
This choice is supported bv thefact that of disk stars have almudauces 0.3DL[nHE]<O15 (Reid&Hawley2t003... but requires that there is no significant coutamunation by other Galactic populations (e.c. thick disk aud halo brown chvarts) iu the observed sample. | This choice is supported by thefact that of disk stars have abundances $-0.3 < [m/H] < 0.15$ \citep{rei00b}, but requires that there is no significant contamination by other Galactic populations (e.g., thick disk and halo brown dwarfs) in the observed sample. |
To convert our fundamental properties to observables. we used the most receit evolutionary calculations roni the Tucson (Burrowsetal.1997). aud Lyou (Baraffeetal.2003). eroups. | To convert our fundamental properties to observables, we used the most recent evolutionary calculations from the Tucson \citep{bur97} and Lyon \citep{bar03}
groups. |
Doth of these models curplov )011-grey: atinosplieres i ivhich condensate opacity is ignored (so-called "COND? nodels: Allard et 22001). arecly consistent with he observed spectra of mid-type M axd 1üid- and late-tvp eT dwarfs (Tsuji.Ohnala.1 1996). | Both of these models employ non-grey atmospheres in which condensate opacity is ignored (so-called “COND” models; Allard et 2001), largely consistent with the observed spectra of mid-type M and mid- and late-type T dwarfs \citep{tsu96}. |
. C1abrieretal.(2000a)} have also derived evolutionary tracks for "DUSTY atiosphere nocdoels. which retaiu condensate material iu their atuospler¢ nore appropriate or wiannucrlate-tvpe A andl L dwarfs (Tsuji.OhuaSiv&Aoki1996). | \citet{cha00b} have also derived evolutionary tracks for “DUSTY” atmosphere models, which retain condensate material in their atmosphere, more appropriate for warmer late-type M and L dwarfs \citep{tsu96}. |
. However. these iithors find 10 difference in the evolukn of unuosity aud Ty¢¢ between he COND aud DUSTY inodes. | However, these authors find $\lesssim$ difference in the evolution of luminosity and $_{eff}$ between the COND and DUSTY models. |
This is a relatively stall deviation give1 the ρουαν larger svstenatic 1ucertaimties arising from the coniplex evolution of condensates iu cool M and L dwarf atinospheres (Ackerman&Marley2001:Burgasseretal.20022:Tsuji2002:Cooper2003) aud current obserational uucertaimties (Bureasser2001:Crzetal.2003). | This is a relatively small deviation given the potentially larger systematic uncertainties arising from the complex evolution of condensates in cool M and L dwarf atmospheres \citep{ack01,me02c,tsu02,coo03} and current observational uncertainties \citep{me01,cru03}. |
. DUSTY evolutionary tracks are herefore iguored 1 this iuvestigation. | DUSTY evolutionary tracks are therefore ignored in this investigation. |
Iu Figure 1.. he evolution of T,;, with time for the two sets of models curploved are compared for nasses 0.001. < AL x04 AL. and ages PM to 10 Cr. | In Figure \ref{fig1}, , the evolution of $_{eff}$ with time for the two sets of models employed are compared for masses 0.001 $\leq$ M $\leq$ 0.1 $_{\sun}$ and ages 1 Myr to 10 Gyr. |
Over much of this parameter space evolutionary racks are consisteif to within 10.. with the Baraffe models predicting sheltly ligher temperatures at a outieular mass ad age for M. < 0.06 AL. and lower temperatures for M. 2 0.08 AL... | Over much of this parameter space evolutionary tracks are consistent to within with the Baraffe models predicting slightly higher temperatures at a particular mass and age for M $<$ 0.06 $_{\sun}$ and lower temperatures for M $>$ 0.08 $_{\sun}$. |
At carly ages (f£<5 Αν) the Burrows nodels are significautly hotter )) for M > 0.06 AL... | At early ages $t \lesssim 5$ Myr) the Burrows models are significantly hotter ) for M $>$ 0.06 $_{\sun}$. |
At later ages (£z5 Gyr). the wo models again ¢leviate significantly }) for 0.06 < ALS 0.08 AL.. with the Daratfe models being th. hotter aud more Iuninous. | At later ages $t \gtrsim 5$ Gyr), the two models again deviate significantly ) for 0.06 $\lesssim$ M $\lesssim$ 0.08 $_{\sun}$, with the Baraffe models being both hotter and more luminous. |
This is due to the higher Uvdrogen Burning Miniuimn Mass (IIBNMO for he Burrows modes, 0.075 versus 0.072 M... | This is due to the higher Hydrogen Burning Minimum Mass (HBMM) for the Burrows models, 0.075 versus 0.072 $_{\sun}$ . |
Finally. the Burrows tracks diverge moresubstantially around he IIBMM.with :v differeuce of 1500 I& between 0.075 aud 0.09 AD. at 10 Corr. as compared to 600 Ix for | Finally, the Burrows tracks diverge moresubstantially around the HBMM,with a difference of 1500 K between 0.075 and 0.09 $_{\sun}$ at 10 Gyr, as compared to 600 K for |
The average offset of longitude is virtually zero; the latitudes have an average offset of -1744. | The average offset of longitude is virtually zero; the latitudes have an average offset of 4. |
This average offset may be due to an underestimate by Hevelius of refraction. | This average offset may be due to an underestimate by Hevelius of refraction. |
The distribution of the total position errors A in lis shown in reff:delta.. | The distribution of the total position errors $\Delta$ in is shown in \\ref{f:delta}. |
This distribution peaks roughly at the value of the width oof the separate distributions in AA, Af, as expected (see explanation in PaperII). | This distribution peaks roughly at the value of the width of the separate distributions in $\Delta\lambda$, $\Delta\beta$, as expected (see explanation in I). |
The number of stars with large position errors is markedly smaller in tthan inKeplerE. | The number of stars with large position errors is markedly smaller in than in. |
. In particular, the number of stars with position errors larger than a degree is 21 (on a total of 1517 identified entries) in aas compared to 47 (on a total of 977 identified entries) inKeplerE. | In particular, the number of stars with position errors larger than a degree is 21 (on a total of 1517 identified entries) in as compared to 47 (on a total of 977 identified entries) in. |
. Similarly, the number of unidentified stars is 16 (of 1533 independent entries) in aand 14 (of 992 independent entries) inKeplerE. | Similarly, the number of unidentified stars is 16 (of 1533 independent entries) in and 14 (of 992 independent entries) in. |
. It may be concluded that the overall accuracy of the star catalogue of Hevelius is better than that of the star catalogue of Brahe/Kepler. | It may be concluded that the overall accuracy of the star catalogue of Hevelius is better than that of the star catalogue of Brahe/Kepler. |
In reff:ksplots we show the cumulative error distributions for each Hevelius magnitude separately, taking magnitudes 1 and 2 together, and limiting the distributions to A«10’. | In \\ref{f:ksplots} we show the cumulative error distributions for each Hevelius magnitude separately, taking magnitudes 1 and 2 together, and limiting the distributions to $\Delta<10$. |
. It is seen that the median error increases slowly but systematically with magnitude. | It is seen that the median error increases slowly but systematically with magnitude. |
In reff:complet we investigate the completeness of aand aas function of magnitude, for three declination ranges. | In \\ref{f:complet} we investigate the completeness of and as function of magnitude, for three declination ranges. |
For this purpose we select from aand oonly those entries which we have identified, and which are not repeat entries, eentries with I=1-4. | For this purpose we select from and only those entries which we have identified, and which are not repeat entries, entries with I=1-4. |
For selecting the Hipparcos stars in the latitude ranges we convert their positions to an equinox halfway between Brahe and Hevelius, 11631.0. | For selecting the Hipparcos stars in the latitude ranges we convert their positions to an equinox halfway between Brahe and Hevelius, 1631.0. |
At magnitudes V «4 there are 348 stars from the with 6>—30°, of which 5 are absent from and 23 from (Cof which 13 with 6« 0°). | At magnitudes $V$$<$ 4 there are 348 stars from the with $\delta>-30^\circ$, of which 5 are absent from and 23 from (of which 13 with $\delta<0^\circ$ ). |
At magnitudes V«5 there are 1138 Hipparcos stars with 6>—30°, of which mmisses 141 stars (89 with 6< 0°) and 3389 (156 with 6< 0°). | At magnitudes $V$$<$ 5 there are 1138 Hipparcos stars with $\delta>-30^\circ$, of which misses 141 stars (89 with $\delta<0^\circ$ ) and 389 (156 with $\delta<0^\circ$ ). |
Finally, about 3500 Hipparcos stars with V<6 have 6>—30°, and of these some 2000 are absent from aand 2500 fromKeplerE,, which is just another way of saying that aand ccontain about 1500 and 1000 stars visible to the naked eye, respectively. | Finally, about 3500 Hipparcos stars with $V$$<$ 6 have $\delta>-30^\circ$, and of these some 2000 are absent from and 2500 from, which is just another way of saying that and contain about 1500 and 1000 stars visible to the naked eye, respectively. |
It may be noted here that the latitude of Gdansk is about 1°55 further south than that of Hven. | It may be noted here that the latitude of Gdansk is about 5 further south than that of Hven. |
Nonetheless, as shown in reff:complet iis more incomplete already at brighter magnitudes, also in the northern parts of the sky. | Nonetheless, as shown in \\ref{f:complet} is more incomplete already at brighter magnitudes, also in the northern parts of the sky. |
How many new stars did Hevelius observe? | How many new stars did Hevelius observe? |
In the manuscript of the catalogue, a note dated 1681 March 31 states that 946 stars of Tycho and 617 new stars were observed (Volkoff et 11971, p.72). | In the manuscript of the catalogue, a note dated 1681 March 31 states that 946 stars of Tycho and 617 new stars were observed (Volkoff et 1971, p.72). |
This gives a total of 1563, very | This gives a total of 1563, very |
between the two sets of sigmoidal loops FP3-FP2 and FP4-FP1 as clearly discernible in EUV images, which results in the J-shaped flare ribbons (alsoseeSchrijveretal.2011). | between the two sets of sigmoidal loops FP3–FP2 and FP4–FP1 as clearly discernible in EUV images, which results in the J-shaped flare ribbons \citep[also see][]{schrijver11b}. |
The reconnected large-scale fields FP3-FP4 could erupt outward to become the halo CME associated with this flare, and the newly formed smaller loops FP1—FP2 lying close to the surface could then account for the enhanced Dj, at the region R. Such a reconnection of two current-carrying loops would also effectively lead the current path to move downward closer to the surface, which can explain the increase of S and 9 (Melrose1997). | The reconnected large-scale fields FP3–FP4 could erupt outward to become the halo CME associated with this flare, and the newly formed smaller loops FP1–FP2 lying close to the surface could then account for the enhanced $B_h$ at the region R. Such a reconnection of two current-carrying loops would also effectively lead the current path to move downward closer to the surface, which can explain the increase of $\tilde{S}$ and $\mathring{S}$ \citep{melrose97}. |
. Alternatively, increase of the magnetic nonpotentiality at and near the surface could result from the newly emerging, sheared magnetic flux (Jingetal.2008),, which could occur after the relaxation of fields above the surface due to the flare energy release. | Alternatively, increase of the magnetic nonpotentiality at and near the surface could result from the newly emerging, sheared magnetic flux \citep{jing08}, which could occur after the relaxation of fields above the surface due to the flare energy release. |
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